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0710 | 0710.2213_arXiv.txt | We investigate gravitino dark matter scenarios in which the primordial $\Lisix$ production is catalyzed by bound-state formation of long-lived negatively charged particles $\champ$ with $\Hefour$. In the constrained minimal supersymmetric Standard Model (CMSSM) with the stau $\stau^-$ as the $\champ$, the observationally inferred bound on the primordial $\Lisix$ abundance allows us to derive a rigid lower limit on the gaugino mass parameter for a standard cosmological history. This limit can have severe implications for supersymmetry searches at the Large Hadron Collider and for the reheating temperature after inflation. | \label{sec:introduction} Big Bang Nucleosynthesis (BBN) is a powerful tool to test physics beyond the Standard Mo\-del. Recently, it has been realized that the presence of heavy long-lived negatively charged particles $\champ$ can have a substantial impact on the primordial light element abundances via bound-state formation% ~\cite{Pospelov:2006sc,Kohri:2006cn,Kaplinghat:2006qr,Cyburt:2006uv,Hamaguchi:2007mp,Bird:2007ge,Kawasaki:2007xb,Jittoh:2007fr,Jedamzik:2007cp}. In particular, when $\champ$ and $\Hefour$ form Coulomb bound states, $(\Hefour\champ)$, too much $\Lisix$ can be produced via the catalyzed BBN (CBBN) reaction~\cite{Pospelov:2006sc} \begin{align} \label{eq:CBBN-reaction} (\Hefour\champ)+\mathrm{D} \rightarrow \Lisix + \champ . \end{align} The formation of $(\Hefour\champ)$ and hence the CBBN production of $\Lisix$ becomes efficient at temperatures $T \sim 10\ \keV$, i.e., at cosmic times $t> 10^3\ \seconds$ at which standard BBN (SBBN) processes are already frozen out. The observationally inferred bound on the primordial $\Lisix$ abundance then restricts severely the $\champ$ abundance at such times. A long-lived $\champ$ may be realized if the gravitino is the lightest supersymmetric particle (LSP). In particular, it is reasonable to consider gravitino LSP scenarios within the constrained minimal supersymmetric Standard Model (CMSSM)~\cite{Ellis:2003dn,Cerdeno:2005eu,Jedamzik:2005dh,Cyburt:2006uv,Pradler:2006hh} in which the gaugino masses, the scalar masses, and the trilinear scalar couplings are parameterized by their respective universal values $\monetwo$, $\mzero$, and $A_0$ at the scale of grand unification $\mgut \simeq 2\times 10^{16}\ \GeV$. Within this framework, the lighter stau $\stau$ is the lightest Standard Model superpartner in a large region of the parameter space and thus a well-motivated candidate for the next-to-lightest supersymmetric particle (NLSP). Since its couplings to the gravitino LSP are suppressed by the (reduced) Planck scale, $\MP=2.4\times 10^{18}\,\GeV$, the stau will typically be long-lived for conserved $\mathrm{R}$-parity% \footnote{For the case of broken R-parity, see, e.g.~\cite{Buchmuller:2007ui}} and thus $\stau^-$ can play the role of $\champ$. In scenarios with conserved $\mathrm{R}$-parity, the gravitino LSP is stable and a promisig dark matter candidate. Gravitinos can be produced efficiently in thermal scattering of particles in the primordial plasma. If the Universe, after inflation, enters the radiation dominated epoch with a high reheating temperature $\TR$, the resulting gravitino density $\Omegatp$ will contribute substantially to the dark matter density $\Omega_{\mathrm{dm}}$~\cite{Bolz:2000fu,Pradler:2006qh,Rychkov:2007uq}. In this work we calculate the amount of $\Lisix$ produced in~(\ref{eq:CBBN-reaction}) by following the treatment of Ref.~\cite{Takayama:2007du}. In particluar, we employ a recent state-of-the-art result for the CBBN reaction cross reaction~\cite{Hamaguchi:2007mp}. The obtained upper limit on the $\champ$ abundance from possible $\Lisix$ overproduction vanishes for sufficiently short $\tau_{\champ}$. This allows us to extract a lower limit on the universal gaugino mass parameter $\monetwo$ within minimal supergravity scenarios where the gravitino is the LSP and the $\champ$ is the $\stau^-$ NLSP.\footnote{In this work we assume a standard cosmological history with a reheating temperature $\TR$ that exceeds the freeze-out temperature $\Tf$ of the $\stau$ NLSP.} This limit leads directly to an upper bound on $\TR$ since $\Omegatp$ cannot exceed the observed dark matter density. The bounds on $\monetwo$ and $\TR$ derived below depend on the gravitino mass but are independent of the CMSSM parameters. Before proceeding, let us comment on the present status of BBN constraints on gravitino dark matter scenarios with a long-lived charged slepton NLSP. In a recent ambitious study~\cite{Jedamzik:2007cp} it is argued that bound-state formation of $\champ$ with protons at $T \sim 1\ \keV$ might well reprocess large fractions of the previously synthesized $\Lisix$. This seems to relax the bound on the $\champ$ abundance for $\tau_{\champ}> 10^6~\seconds$. However, at present, the uncertainties in the relevant nuclear reaction rates in~\cite{Jedamzik:2007cp} make it difficult to decide whether a new cosmologically allowed region will open up. In this work we assume that this is not the case, in particular, since the $^3$He/D constraint on electromagnetic energy release~\cite{Sigl:1995kk} becomes severe in this region and excludes stau lifetimes $\tau_{\stau}\gtrsim10^6~\seconds$~\cite{Cerdeno:2005eu,Cyburt:2006uv,Kawasaki:2007xb,Jedamzik:2007cp}. Then only the constraint from hadronic energy release on D~\cite{Kawasaki:2004qu,Feng:2004mt,Cerdeno:2005eu,Jedamzik:2006xz,Steffen:2006hw} can be slightly more severe than the one from catalyzed $^6$Li production~\cite{Cyburt:2006uv,Steffen:2006wx,Pradler:2006hh,Kawasaki:2007xb}. We neglect the D constraint in this work since it can only tighten the bounds on $\monetwo$ and $\TR$ as can be seen, e.g., in Figs.~4~(b--d) and~5 of Ref.~\cite{Pradler:2006hh}. For deriving conservative bounds on $\monetwo$ and $\TR$, it is thus sufficient to consider the CBBN reaction~(\ref{eq:CBBN-reaction}) exclusively. | \label{sec:conclusion} We have considered the catalysis of $\Lisix$ production in CMSSM scenarios with the gravitino LSP and the stau NLSP. Within a standard cosmological history, the calculated $\Lisix$ abundance drops below the observational limit on primordial $\Lisix$ for $\tau_{\stau} \lesssim 5\times 10^3\,\seconds$. Taken at face value, we find that this constraint translates into a lower limit $\monetwo \ge 0.9\, \TeV ( \mgr / 10\, \GeV )^{2/5}$ in the entire natural region of the CMSSM parameter space. This implies a conservative upper bound $\TR\lesssim 4.9\times 10^7\GeV( \mgr / 10\, \GeV )^{1/5}$. The bounds on $\monetwo$ and $\TR$ not only confirm our previous findings~\cite{Pradler:2006hh} but are also independent of the particular values of the CMSSM parameters for the considered $\stau$ NLSP abundances. | 7 | 10 | 0710.2213 |
0710 | 0710.2814_arXiv.txt | CO isotopes are able to probe the different components in protostellar clouds. These components, core, envelope and outflow have distinct physical conditions and sometimes more than one component contributes to the observed line profile. In this study we determine how CO isotope abundances are altered by the physical conditions in the different components. We use a 3D molecular line transport code to simulate the emission of four CO isotopomers, $^{12}$CO $J=2\rightarrow1$, $^{13}$CO $J=2\rightarrow1$, C$^{18}$O $J=2\rightarrow1$ and C$^{17}$O $J=2\rightarrow1$ from the Class 0/1 object L483, which contains a cold quiescent core, an infalling envelope and a clear outflow. Our models replicate JCMT (James Clerk Maxwell Telescope) line observations with the inclusion of freeze-out, a density profile and infall. Our model profiles of $^{12}$CO and $^{13}$CO have a large linewidth due to a high velocity jet. These profiles replicate the process of more abundant material being susceptible to a jet. C$^{18}$O and C$^{17}$O do not display such a large linewidth as they trace denser quiescent material deep in the cloud. | Molecules, particularly CO, are used as tracers of H$_2$ and thus of gas density in cold dark clouds. CO is highly abundant with a low critical density and typically exhibits strong optical depth effects in cold dark molecular clouds. The four most common CO isotopes differ in abundance by as much as three orders of magnitude. Thus they become optically thick at different column densities. Taken together, observations of CO isotopes can trace the gas density in all the main components of cold, dark clouds: the intermediate optical depth envelope, the high optical depth core, and the optically thin bipolar outflows. However, we know from observational and theoretical studies that the abundance of CO depends on conditions in the clouds such as, shock heating \citep{van_dishoeck_95,nisini_07}, U.V. excitation \citep*{goldsmith_07}, freeze-out \citep*{lee_04} and varies from place to place. Therefore we cannot use CO as an H$_2$ tracer without understanding its chemical variation. To better understand CO variations in cold dark clouds, we observed and modeled one particular cloud, Lynds 483. We chose L483 as a prototype for a study of CO abundances because it is a well studied nearby ($\sim 200~{\rm pc}$) molecular cloud. It contains an IRAS source 18148-0440 that is in transition between a Class 0 and Class 1 object \citep{tafalla.et.al00}. It exhibits an infalling envelope \citep{park.et.al99,park.et.al00,tafalla.et.al00} and a slow bipolar molecular outflow \citep{fuller.et.al95,buckle.et.al99} yet the core and envelope are still cold and dense \citep{ladd.et.91,fuller.et.92,fuller&wootten00}. Thus many of the physical properties and kinematic features that are present in either less or more evolved clouds are all present in L483. Our aim is to combine these components into a single model for L483 to strongly constrain the structure and dynamics of the system and hence then to infer the CO abundance throughout the cloud. \begin{figure*} \centering \includegraphics[width=180mm]{fig1} \caption{An example line profile of each of the four isotopes used in this study where the offsets with respect to IRAS 18140-0440 are indicated in the top right corner of each panel. The $^{12}$CO y-axis is in units of $T_{r}$ and $^{13}$CO, C$^{18}$O and C$^{17}$O are in units of $T_{mb}$ } \label{overview} \end{figure*} We obtained emission line profile data from L483 for transitions of the four most common CO isotopes (section ~\ref{obs}). In addition, we used an archival dust continuum emission map of L483 as an unbiased mass tracer to be compared with the CO data. We used a radiative transfer model to analyse the data and to produce synthetic spectra to be compared in detail with individual observations (section ~\ref{models}). Estimates for density, temperature and a physical model of L483 were constrained with help from the observational literature. The abundances of the CO isotopes were then varied in the model to give a good match with the observed line profiles. This yields the abundance variation throughout L483 as well as self-consistent temperatures, densities and velocities and abundance ratios. | We modelled L483 in four isotopomers of CO in order of decreasing abundance $^{12}$CO, $^{13}$CO, C$^{18}$O and C$^{17}$O. Each species delineates a different region and therefore we get a clearer picture of cloud dynamics. C$^{18}$O and C$^{17}$O are optically thin and are important tracers of the denser regions of star forming clouds. Optically thick species such as $^{12}$CO and $^{13}$CO are more abundant and more susceptible to the jet motion. They trace regions farther out from the centre of the cloud and have line profiles showing a large line width consistent with material exposed to a high velocity jet. Using a radiative transfer code and a dynamical model for L483 we were able to self-consistently calculate for the first time the abundance of the CO isotopes in the different regions of such a cloud. Our principal finding is that the CO line profiles in L483 are well fitted with a self-consistent envelope plus boundary layer model and that the CO abundances increase substantially in this boundary layer. The most likely reason for this is that molecular ices on dust grains are heated and released back into the gas phase in the boundary layer. A constant abundance model was found to overestimate the abundance towards the centre of the cloud and only freeze-out of material towards the centre was able to produce modeled profiles consistent with observations. Our tanh geometry is chosen because it matches the observed morphology seen in other protostellar outflows \citep{tafalla.et.al00}. Other geometries are possible for the outflow e.g. conical, cylindrical outflow but it is unlikely that it would have a substantial effect on our results because such a detailed treatment of the boundary layer components is difficult to achieve until there is sub-arcsecond resolution resolution. We emphasize that our results provide an abundance enhancement measurement rather than proving an exact mechanism by which the CO is enhanced, e.g. chemical reactions or dissociation. A more detailed treatment would involve a full dark cloud and gas-grain chemistry whilst accounting for localized shock heating. The most enhanced species in our study, by a factor of $\sim 30$, is the $^{13}$CO material. The exothermic reaction leading to the creation of $^{13}$CO \citep{duley_williams} is shown below \begin{equation} \; \; \; \; \; ^{13}{\rm C}^{+} + ^{12}{\rm CO} \rightleftharpoons ^{12}{\rm C}^{+} + ^{13}{\rm CO} + \Delta{\rm E} \label{13co_creation} \end{equation} where the zero-point energy difference $\Delta{\rm E}$ is equivalent to a temperature $\Delta{\rm E}/{\rm k}$ of 35~K. This mechanism may be the source for the enhanced abundance observed from our modeling. The enhanced abundance seen in the boundary layer effect may also be detectable in other molecules. \citet{park.et.al00} used interferometric observations of HCO$^{+}$ and observed anti-infall profiles close to the centre of the cloud. They concluded the HCO$^{+}$ was tracing the outlying regions of the outflow, i.e. a region between the envelope and the jet. The reason the HCO$^{+}$ emission is predominately seen here rather than in the more extensive envelope is also likely due to an enhancement of HCO$^{+}$ caused by the shock-heated release of icy grain mantles followed by chemical reaction. CO and H$_2$O are liberated into the gas phase and the shock-induced radiation field then can photodissociate CO to C$^+$. This then reacts with the H$_2$O to form HCO$^{+}$. Such a model was successfully used to explain the enhancement of HCO$^{+}$ commonly seen at the bases of molecular outflows \citep{rawlings.et.al00,rawlings.et.al04}. The results in this paper demonstrate that a combination of datasets with several lines and transitions coupled with a 3D molecular line transport code is a powerful way to determine the properties of dense star forming cores. | 7 | 10 | 0710.2814 |
0710 | 0710.5119_arXiv.txt | {The excess above 1 GeV in the energy spectrum of the diffuse Galactic gamma radiation, measured with the EGRET experiment, can be interpreted as the annihilation of Dark Matter (DM) particles. The DM is distributed in a halo around the Milky Way. Considering the directionality of the gamma ray flux it is possible to determine the halo profile. The DM within the halo has a smooth and a clumpy component. These components can have different profiles as suggested by N-body simulations and the data is indeed compatible with a NFW profile for the diffuse component and a cored profile for the clumpy component. These DM clumps can be partly destroyed by tidal forces from interactions with stars and the gravitational potential of the Galactic disc. This effect mainly decreases the annihilation signal from the Galactic centre (GC). In this paper constraints on the different profiles and the survival probability of the clumps are discussed. \PACS{ {95.35.+d}{Dark Matter} \and {98.35.Gi}{Galactic halo} } % } % | \label{seq:intro} From WMAP measurements of the temperature aniso- tropies in the Cosmic Microwave Background (CMB) in combination with data on the Hubble expansion and the density fluctuation in the Universe \cite{RefSpergel} we gather that Cold Dark Matter (CDM) makes up 23\% of the energy of the Universe. The nature of the Dark Matter (DM) is unknown, but one of the most promising particles is the "weakly interacting massive particle" (WIMP). Assuming that WIMPS are Majorana particles they can annihilate each other and produce a large amount of secondary particles. For the determination of the density distribution, the so-called DM halo profile of the WIMP particles, the gamma radia\-tion is most important because it is not influenced by the magnetic field of the galaxy and points back directly to its source.\\ The observation of the diffuse Galactic gamma radiation of the Milky Way with EGRET showed an excess above 1 GeV in the photon energy spectrum. This excess is different for various sky directions and can be interpreted as a WIMP annihilation signal \cite{RefSander}. Therefore it can be used to determine the DM halo profile.\\ In section \ref{seq:halo}, we will describe the mechanism of the determination of the DM halo profile from the EGRET excess and explain how to calculate the DM annihilation flux of gamma rays. Then, after differentiating between diffuse and clumpy DM, we will introduce a survival probability for DM clumps as well as a ringlike substructure of DM within the Galactic disc. | CDM probably consists of WIMPS which are heavy and slow particles. If these particles are Majorana particles they can annihilate each other and produce Galactic gamma radiation which can be used to determine the density profile of the DM. In this analysis the directionality of the DM annihilation flux measured with EGRET was used to find a possible DM halo profile. After dividing the Galactic gamma ray flux into 4 latitude and 45 longitude regions the background and the DM annihilation signal were fitted to the data for each of the 180 bins. The DM annihilation flux is dominated by the annihilation flux of the clumpy DM which is proportional to $\rho_{\chi, clump}$, not $\rho^2_{\chi, clump}$. While the diffuse DM component has a cuspy NFW profile a shallower cored distribution was obtained for the clumpy component. The DM annihilation flux is dominated by the clumpy DM component, but the clumpy component yields a mass below the required mass $> 10^{12}$ solar masses \cite{RefBattaglia}. However, if combined with the diffuse cuspy NFW profile, both the EGRET data and the mass constraint can be fulfilled. In order to take the tidal disruption of DM clumps in the vicinity of stars into account a survival probability for clumps was introduced. Most of the clumps are expected to be destroyed near the Galactic centre, although a steep cusp may survive. The strong signal observed from the Galactic centre yielded a survival probability at the centre of $P(0)=0.7$. This means that the DM clumps are not completely destroyed, which is in good agreement with more detailed calculations in Ref. \cite{RefDokuchaev2}. In addition to the DM halo profiles two ringlike substructure were required at radii of 4 and 13 kpc. The halo and ring parameters were obtained by minimizing a $\chi^2$ function comparing the flux of the excess from the various sky directions with the line-of-sight integral in the halo. Figure \ref{fig:longitudes} shows that the halo model fits the measured data very well.\\ In summary, the EGRET excess of diffuse Galactic gamma rays is in good agreement with the expectations of a cored clumpy halo component plus a cuspy diffuse one. The ringlike substructure, expected from the tidal disruption of the nearby Canis Major dwarf galaxy, is clearly seen and its heavy mass above $10^{10}$ solar masses as obtained from the EGRET data, has been recently confirmed by the reduced gas flaring in this region. | 7 | 10 | 0710.5119 |
0710 | 0710.5269_arXiv.txt | We investigate the full $5D$ dynamics of general braneworld models. Without making any further assumptions we show that cyclic behavior can arise naturally in a fraction of physically accepted solutions. The model does not require brane collisions, which in the stationary case remain fixed, and cyclicity takes place on the branes. We indicate that the cosmological constants play the central role for the realization of cyclic solutions and we show that its extremely small value on the observable universe makes the period of the cycles and the maximum scale factor astronomically large. | The last decade proves to be really exciting for cosmology. Observational data indicated, among other very interesting results, that the expansion of the universe is accelerated \cite{observ}. At the same time the braneworld scenario appeared in the literature \cite{Horawa,RS99}. Though the exciting idea that we live in a fundamentally higher-dimensional spacetime which is greatly curved by vacuum energy was older \cite{Rubakov83}, the new class of ``warped" geometries offered a simple way of localizing the low energy gravitons on the brane. In this novel background the old idea of a cyclic Universe was reheated. Started as ekpyrotic \cite{Khoury.best,Khoury.rebound}, enriched to ekpyrotic/cyclic \cite{ekpyrotic1,Khoury.rebound,Turok.simplified,Turok.sing,ekpyrotic.pert,cyclic.clifton,chargeBH} and recently to new ekpyrotic \cite{ekpyrotic3,ekpyrotic4,ekpyrotic2,ekpyrotic.fields}, the new paradigm tries to be established as an alternative to standard cosmology. According to its basic contents, our universe experiences an infinite or extremely large number of cycles, each one consisting of a hot bang phase, a phase of accelerated expansion, a phase of slow-ekpyrotic contraction and a bounce-bang that triggers the next cycle. Starting with a simplified notional framework (infinite and not ``created" time) cyclic cosmology have many advantages. It successfully faces the homogeneity, isotropy, topological and flatness problems, it handles the issue of initial conditions, it incorporates the dark energy and transforms it to an important factor, and it provides the mechanism of the generation of cosmic perturbations and of structure formation. However, there are some key issues that do not have a consistent and efficient approach so far, despite the great progress. These are the settlement of the singularity, although temperature and density remain finite, the entropy evolution, and the fate of the perturbations through the bounce. Through this research, cyclic scenarios have become more complicated, by the insertion of more complex potentials, of more branes \cite{Khoury.best}, of the mechanism of ghost condensation \cite{ghcond,ekpyrotic3}, of more scalar fields \cite{ekpyrotic.fields} and of procedures which cancel the tachyonic instabilities \cite{ekpyrotic4}. Most of the works on cyclic cosmology involve, initially or at some stage, the transition to effective $4D$ equations. However, as it was mentioned in \cite{Linde01,4Dbreakdown}, such a procedure does not lead to reliable results since one cannot return to the $5D$ description self-consistently. Furthermore, the old $4D$-singularity problem (of both Big Bang and traditional cyclic universes), has been replaced by a new one (equally annoying) concerning the singularity of extra dimension(s). This later case is accompanied by the brane collision phenomenon, which seems to be a basic constituent of the ekpyrotic scenario. In this work we desire to investigate the full $5D$ dynamics of general braneworld models and examine if a cyclic behavior is possible. This is an essential procedure in order to consistently confront the arguments of the authors of \cite{Linde01}, which claim that cyclic behavior cannot arise from a complete $5D$ description, and our study must not include any additional assumptions or fine tunings in order to remain general and therefore convincing. Secondly, we are interested to explore if a cyclic behavior of $5D$ dynamics is necessarily related to brane collisions. This work is organized as follows: In section \ref{model} we present the $5D$ braneworld model and we derive the equations of motion. In section \ref{analyt} we provide analytical solutions for two simplified stationary solution subclasses, while in \ref{numer} we investigate numerically the full stationary dynamics. Finally, in section \ref{discussion} we discuss the physical implications of our analysis and we summarize the obtained results. | \label{discussion} In the aforementioned analysis we considered general braneworld models characterized by the action (\ref{action}), the conformal metric (\ref{metric}), and the general potentials (\ref{bulkpot}) and (\ref{branepot}). Performing both analytical and numerical calculations we showed that the full $5D$ dynamics allows for stationary solutions corresponding to oscillatory scale factor of the physical brane and therefore to cyclic universes. In statistical terms cyclicity corresponds to $\approx 4\%$ of the physical solutions. Our investigation is completely $5D$, cyclic behavior arises naturally and is induced on the brane by the full dynamics, and it is not a result of a modified $4D$ dynamics, with fine-tuned parameters or specific assumptions in the Friedmann equation. Furthermore, we do not use an explicit brane state equation, considering just the bulk scalar field (the decays and interactions of which will eventually fill the physical brane with the conventional content \cite{apostol}). As we mentioned in the introduction this full $5D$ approach is necessary in order to confront the arguments of the authors of \cite{Linde01}. Indeed, their allegations that one cannot transit to an effective $4D$ theory (integrating the action over $y$), solve the equations there and then return naively to the $5D$ description (adding time-dependence by hand), are correct. Doing so, the results are not self-consistent (especially the boundary conditions are not satisfied) and the authors of \cite{Linde01} use this fact as a central argument against the cyclic scenario. However, our consistent $5D$ analysis reveals that cyclic behavior is possible. Another important feature of the present study is that cyclic universes do not require brane collisions. Thus, we avoid the known problems concerning such a description, which force ekpyrotic model to successively more complicated versions. On the contrary, the branes do not move at all and the system is stable (stationary solutions are a stable fixed point \cite{brcod,Tetradis01}). Furthermore, in our model, expansion and contraction take place in the 3+1 branes, and in all 3+1 slices in general, while the fifth dimension remains unaffected. The $4$ spacial dimensions shrink periodically to an $1D$ string and re-expand. This is in a radical contrast with the cyclic models with extra dimensions, where the extra dimension is the one that gets contracted (the fifth in \cite{ekpyrotic1} or the eleventh in \cite{Khoury.rebound}). Cyclicity seems to re-obtain its ``physical" meaning. Our $5D$ investigation is general and does not involve extra assumptions, fine-tunings or specific potential forms. We result to periodic, cyclic, homogenous and isotropic universes, where the scale factor changes smoothly from expanding to contracting. An observer on the physical brane feels successively accelerated expansion, decelerated expansion, turnabout, accelerated contraction, decelerated contraction, bounce e.t.c, and a promising signature of the cyclic behavior would be the measure of the varying rate of the Hubble constant. The cycles period, given in (\ref{period}), can be arbitrary, depending on $B(0)$, i.e. on the value of the warp factor on the physical brane ($\theta$ is bounded from above and therefore cannot act as a period-decreasing factor). A very interesting conclusion comes from the insertion of observational results in our model, which was not made above in order to remain as general as possible. Explicitly, if we use the fact that the cosmological constant of our Universe is extremely small ($\approx{\cal O}(10^{-47})\ \text{GeV}^4$), and assuming a reasonable $M_5$ value of ${\cal O}(10^{19})$ GeV, the first two boundary conditions in relation (\ref{junctions}) provide in general a huge value for $e^{B(0)}$ ($\approx{\cal O}(10^{45})$). This is in consistency with the scaling transformation of \cite{brcod}, which allows us, in a solution, to scale the parameters by $e^{-S}$ and add to the warp factor the constant $S$, and acquire another solution. Therefore, the extremely small cosmological constant of the observable universe leads the cyclicity period to be around $T\approx{\cal O}(10^{13})$ years and the maximum scale factor value, given by (\ref{minmax}), to be $a_{max}\approx {\cal O}(10^{28})$ m (where the decimal exponents in these rough estimations can vary by 1 or 2, depending on $B'(0)$ and $\theta$ values). Luckily enough, the smallness of the cosmological constant excludes oscillatory models with small periods in astronomical terms. In more foundational words, the reason that made the cosmological constant that small, is the same that makes the cycle period and the size of the Universe that large. In this work we have been restricted to stationary solutions, where the subclass of them that possesses $H^2<0$ corresponds to eternal cyclic behavior with constant period. Numerical investigation of the full dynamics seem to consist of such stationary solutions and the transitions between them \cite{brcod,Tetradis01}. In such transitions $H_0^2$ on the physical brane can chance sign, leading to a form of ``chaotic cyclicity", where large intervals of (non-periodic in general) oscillatory behavior could be followed by large intervals of conventional evolution and vice versa. In this case, an initial Big Bang and/or a final Big Rip or Big Crunch (in conventional terms) could be possible. Another interesting possibility would be the exploration of our model with cosmological constants being piecewise constant functions of time, reflecting cosmological phase transitions, which could also lead to chaotic cyclicity. Note however that numerical confirmation of such behaviors is very hard due to the small probability of cyclic stationary solutions ($\approx 10^{-2}\%$ as we have already mentioned). These subjects are under investigation. In order for a model to serve as a description of nature, it has to explain the basic physical key issues. Especially for cyclic cosmology, amongst others these are the entropy evolution and, probably the most pressing issue, that of a fuller understanding of the bounce and the handling of the singularity. Our model provides a consistent background for cyclicity and it reveals how such a behavior arises from the full $5D$ dynamics. However, since braneworlds and brane cosmology in general arise as limits of a multi-dimensional theory unknown up to now, the $5D$ results have a phenomenological character and must be considered from this point of view. Definitely, a complete explanation and apprehension, and a successful confrontation of the aforementioned subjects, can only come through a higher-dimensional, fundamental theory of nature. For the moment we have to rely on the relevant research on cyclic cosmology, linearized gravity, M-theory and strings, which has improved our knowledge on these issues. These results can be embodied in our analysis. The most hopeful effort is the use of quantum fluctuations in order to tame the singularity, which effectively is translated into a modification of gravity by the scalar field \cite{Bojowald01,ekpyrotic3,ekpyrotic4}. Alternatively, using loop quantum gravity we could modify non-perturbatively the dynamical equations leading to a singularity resolution as in \cite{Maartens}. Concerning the entropy, we could include the relevant discussion in our investigation. The argument of the authors of \cite{Turok.sing,Turok.simplified} about maximum amount of entropy possible in de Sitter spacetime, may lead our model to have a maximum cycle number between $10^{20}$ and $10^{30}$. However, the idea of the causal patch \cite{Turok.sing} is probably the best way of handling the entropy problem so forth, and there are some interesting recent works on the subject which give a boost on cyclic cosmology \cite{entropy}. Let us close this discussion section with some comments on the role of the brane tensions and of the bulk cosmological constant in our model. As can be numerically confirmed, setting them to zero makes it almost impossible to satisfy the boundary conditions obtaining $H^2<0$ and singularity absence in [0,1] (this can be achieved only through a careful fine-tuning since our random choice procedure gives an one-digit number of such solutions in $10^6$ parameter multiplets). On the other hand, as we showed in \ref{B}, in the case where $\Lambda$, $\lambda_0$ and $\lambda_1$ are the only non-zero parameters, an $\approx10^{-1}\%$ of the random parameter choices, or $\approx6\%$ of the solutions, correspond to $H^2<0$. In mathematical terms, $\Lambda$, $\lambda_0$ and $\lambda_1$ are requisite in order to acquire a solution with $H^2<0$ in the full dynamics, in a natural and not in a fine-tuning way. In terms of physics, it is the dark energy that lies in the background of the oscillatory mechanism and allows for cyclicity to realize. Adding the fact that it determines the cycles period and the maximum scale factor value, we conclude that dark energy is crucial in the described model. This brings it closer to the ekpyrotic paradigm of the literature. In this work we examine general braneworld models and we show that cyclic behavior can naturally arise from the full $5D$ dynamics. One important feature is that brane collisions are not required, on the contrary the branes remain stable, and the cyclicity takes place on the $4D$ geometry not on the extra dimension. Another significant result is that the smallness of the cosmological constant of the observable universe pushes the cyclic period and the scale factor to astronomical large values, an essential requirement for the establishment of cyclic cosmology as a realistic alternative paradigm. Furthermore, we indicate the possibility of a ``chaotic cyclicity", that is extremely large, non-periodic, cyclic intervals followed by extremely large intervals of conventional evolution and vice versa. After these, the model shares both the advantages and disadvantages of cyclic cosmology.\\ \paragraph*{{\bf{Acknowledgements:}}} The author acknowledges partial financial support through the research program ``Pythagoras'' of the EPEAEK II (European Union and the Greek Ministry of Education). | 7 | 10 | 0710.5269 |
0710 | 0710.5096_arXiv.txt | We investigate various galaxy occupation statistics of dark matter halos using a large galaxy group catalogue constructed from the Sloan Digital Sky Survey Data Release 4 (SDSS DR4) with an adaptive halo-based group finder. The conditional luminosity function (CLF), which describes the luminosity distribution of galaxies in halos of a given mass, is measured separately for all, red and blue galaxies, as well as in terms of central and satellite galaxies. The CLFs for central and satellite galaxies can be well modelled with a log-normal distribution and a modified Schechter form, respectively. About 85\% of the central galaxies and about 80\% of the satellite galaxies in halos with masses $M_h\ga 10^{14}\msunh$ are red galaxies. These numbers decrease to 50\% and 40\%, respectively, in halos with $M_h \sim 10^{12}\msunh$. For halos of a given mass, the distribution of the luminosities of central galaxies, $L_c$, has a dispersion of about $0.15$ dex. The mean luminosity (stellar mass) of the central galaxies scales with halo mass as $L_c \propto M_h^{0.17}$ ($M_{*,c} \propto M_h^{0.22}$) for halos with masses $M\gg 10^{12.5}\msunh$, and both relations are significantly steeper for less massive halos. We also measure the luminosity (stellar mass) gap between the first and second brightest (most massive) member galaxies, $\log L_1 - \log L_2$ ($\log M_{*,1}-\log M_{*,2}$). These gap statistics, especially in halos with $M_h \la 10^{14.0} \msunh$, indicate that the luminosities of central galaxies are clearly distinct from those of their satellites. The fraction of fossil groups, defined as those groups with $\log L_1 - \log L_2\ge 0.8$, ranges from $\sim 2.5\%$ for groups with $M_h\sim 10^{14}\msunh$ to 18-60\% for groups with $M_h\sim 10^{13}\msunh$. The number distribution of satellite galaxies in groups of a given mass follows a Poisson distribution, in agreement with the occupation statistics of dark matter sub-halos. This provides strong support for the standard lore that satellite galaxies reside in sub-halos. Finally, we measure the fraction of satellites, which changes from $\sim 5.0\%$ for galaxies with $\rmag\sim -22.0$ to $\sim40\%$ for galaxies with $\rmag\sim -17.0$. | In recent years, the halo occupation distribution and conditional luminosity function have become powerful statistical measures to probe the link between galaxies and their hosting dark matter halos. Although these statistical measures themselves do not give physical explanations of how galaxies form and evolve, they provide important constraints on various physical processes that govern the formation and evolution of galaxies, such as gravitational instability, gas cooling, star formation, merging, tidal stripping and heating, and a variety of feedback processes. In particular, they constrain how their efficiencies scale with halo mass. The halo occupation distribution (hereafter HOD), $P(N \vert M)$, which gives the probability of finding $N$ galaxies (with some specified properties) in a halo of mass $M$, has been extensively used to study the galaxy distribution in dark matter halos and galaxy clustering on large scales (e.g. Jing, Mo \& B\"orner 1998; Peacock \& Smith 2000; Seljak 2000; Scoccimarro \etal 2001; Jing, B\"orner \& Suto 2002; Berlind \& Weinberg 2002; Bullock, Wechsler \& Somerville 2002; Scranton 2002; Kang \etal 2002; Marinoni \& Hudson 2002; Zheng \etal 2002; Magliocchetti \& Porciani 2003; Berlind \etal 2003; Zehavi \etal 2004, 2005; Zheng \etal 2005; Tinker \etal 2005). The conditional luminosity function (CLF), $\Phi(L \vert M) {\rm d}L$, which refines the HOD statistic by considering the average number of galaxies with luminosity $L \pm {\rm d}L/2$ that reside in a halo of mass $M$, has also been extensively investigated (Yang, Mo \& van den Bosch 2003; van den Bosch, Yang \& Mo 2003; Vale \& Ostriker 2004, 2006; Cooray 2006; van den Bosch et al. 2007a) and has been applied to various redshift surveys, such as the 2dFGRS, the Sloan Digital Sky Survey (SDSS) and DEEP2 (e.g. Yan, Madgwick \& White 2003; Yang \etal 2004; Mo et al. 2004; Wang \etal 2004; Zehavi \etal 2005; Yan, White \& Coil 2004). These investigations demonstrate that the halo occupation statistics are very useful in establishing and describing the connection between galaxies and dark matter halos. Furthermore, they also indicate that the galaxy/dark halo connection can provide important constraints on cosmology (e.g.,van den Bosch, Mo \& Yang 2003; Zheng \& Weinberg 2007). Finally, the HOD/CLF framework also allows one to split the galaxy population in centrals and satellites, and to describe their properties separately (e.g. Cooray 2005; White \etal 2007; Zheng \etal 2007). As has been pointed out in Yang et al. (2005c; hereafter Y05c), a shortcoming of the HOD/CLF models is that the results are not completely model independent. Typically, assumptions have to be made regarding the functional form of either $P(N \vert M)$ or $\Phi(L \vert M)$. Moreover, in all HOD/CLF studies to date, the occupation distributions have been determined in an indirect way: the free parameters of the assumed functional form are constrained using {\it statistical} data on the abundance and clustering properties of the galaxy population. One may hope to circumvent this shortcoming by directly measure the dark matter distribution around galaxies. Such measurements can in principle be obtained through gravitational lensing and X-ray observations. However, both methods are hampered by requirements on the data quality and uncertainties related to the interpretation of the data. For instance, weak lensing measurements, which requires high-quality imaging, typically needs to resort to the stacking of many lens galaxies in order to get a detectable signal, but this stacking severely complicates the interpretation in terms of the halo masses of the lens galaxies. In the case of X-ray observations, robust constraints can only be obtained for massive clusters, but even here the interpretation of the data can be complicated due to the presence of substructure and deviations from hydrostatic equilibrium. An alternative method to directly probe the galaxy - dark halo connection is to use galaxy groups as a representation of dark matter halos and to study how the galaxy population changes with the properties of the groups (e.g., Y05c; Zandivarez et al. 2006; Robotham et al. 2006; Hansen \etal 2007). Recently, we have constructed a large galaxy group catalogue based on the Sloan Digital Sky Survey Data Release 4 (SDSS DR4), using an adaptive halo-based group finder (Yang \etal 2007; Paper I; Y07 hereafter). Detailed tests with mock galaxy catalogues have shown that this group finder is very successful in associating galaxies according to their common dark matter halos. In particular, the group finder performs reliably not only for rich systems, but also for poor systems, including isolated central galaxies in low mass halos. This makes it possible to study the galaxy-halo connection for systems covering a large dynamic range in masses. Various observational selection effects have been taken into account, especially the survey edge effects and fiber collisions. The halo masses for the groups are estimated according to the abundance match, using the characteristic group luminosity and stellar masses (see \S\ref{sec_data} below). According to tests with mock galaxy catalogues, the halo masses are estimated with a standard deviation of about 0.3 dex. With these well-defined galaxy group catalogues, one can not only study the properties of galaxies in different groups (e.g. Y05c; Yang \etal 2005d; Collister \& Lahav 2005; van den Bosch \etal 2005; Robotham \etal 2006; Zandivarez \etal 2006; Weinmann \etal 2006a,b; van den Bosch \etal 2007b; McIntosh \etal 2007), but also probe how dark matter halos trace the large-scale structure of the Universe (e.g. Yang \etal 2005b, 2006; Coil \etal 2006; Berlind \etal 2007; Wang et al. 2007 in preparation). In this paper, which is the second in a series, we use the SDSS DR4 group catalogue to probe various occupation statistics and measure the CLFs for different populations of galaxies. This paper is organized as follows: In Section~\ref{sec_data} we describe the data (galaxy and group catalogues) used in this paper. Section~\ref{sec_CLFs} presents our measurement of the CLFs for all, red and blue galaxies. Sections~\ref{sec_central}, ~\ref{sec_HOD} and ~\ref{sec_satfrac} describe the properties of central galaxies, the halo occupation statistics and the fraction of satellite galaxies, respectively. Finally, we summarize our results in Section~\ref{sec_summary}. Throughout this paper, we use a $\Lambda$CDM `concordance' cosmology whose parameters are consistent with the three-year data release of the WMAP mission: $\Omega_m = 0.238$, $\Omega_{\Lambda}=0.762$, $n_s=0.951$, $h=0.73$ and $\sigma_8=0.75$ (Spergel et al. 2007). \begin{figure} \plotone{f1.eps} \caption{The color-magnitude relation for galaxies in our group sample. The open circles indicate the Gaussian peaks of the bi-normal distribution of galaxies in each luminosity bin. The solid dots indicate the corresponding averages of the two Gaussian peaks. The solid line is the best-fit quadratic relation to these averages (see eq.~[\ref{quadfit}]), which we use to split the galaxies into red and blue population (color-coded accordingly).} \label{fig:data_color} \end{figure} \begin{figure*} \plotone{f2.eps} \caption{The conditional luminosity functions (CLFs) of galaxies in groups of different mass bins. Symbols correspond to the CLFs obtained using $M_L$ as halo mass (estimated according to the ranking of the characteristic group luminosity), with solid and open circles indicating the contributions from central and satellite galaxies, respectively. The errorbars reflect the 1-$\sigma$ scatter obtained from 200 bootstrap samples. The solid lines indicate the related best-fit parameterizations using equation~[\ref{eq:CLF_fit}]. For comparison, we also show, with dashed lines, the CLFs obtained using $M_S$ as halo mass (estimated according to the ranking of the group's characteristic stellar mass). Results shown in this plot are obtained from Sample II. } \label{fig:CLF} \end{figure*} \begin{figure*} \plotone{f3.eps} \caption{Similar to Fig.~\ref{fig:CLF}, but here we show the CLFs for red (dashed lines) and blue (dotted lines) galaxies, for groups with halo masses $M_L$. The solid lines indicate the best-fit parameterizations for the CLFs of red galaxies. In both cases the central and satellite components of the CLFs are indicated separately. For clarity, the errorbars, again obtained using 200 bootstrap samples, are only shown for the red galaxies. } \label{fig:CLF_color} \end{figure*} \begin{figure*} \plotone{f4.eps} \caption{The best fit parameters ($\phi_s^{\star}$ upper row, $\alpha_s^{\star}$ second row, $L_c$ third row, and $\sigma_c$ bottom row) to the CLFs shown in Figs.~\ref{fig:CLF} and ~\ref{fig:CLF_color}, as functions of halo mass. Panels on the left, in the middle, and on the right show results for all, red, and blue galaxies, respectively. Since we have two different halo mass estimators ($M_L$ and $M_S$) and two main group samples (II and III), we have obtained CLFs for four different combinations of sample and group mass estimator. The results for all four combinations are shown using different symbols and line-styles, as indicated. The errorbars in the first two and last rows indicate the 1-$\sigma$ variances obtained from our 200 bootstrap samples. In the third row of panels, however, the errorbars correspond to the log-normal scatter, $\sigma_c$, shown in the bottom row of panels. For clarity the errorbars are only shown for the `$M_L$-Sample II' case, but they are very similar for the other four cases. } \label{fig:fit_CLF} \end{figure*} | \label{sec_summary} Using a large galaxy group catalogue constructed from the SDSS Data Release 4 (DR4) by Y07, we have investigated various halo occupation statistics of galaxies. In particular, we have split the galaxy population in red and blue galaxies, and in centrals and satellites, and determined the conditional luminosity functions of these varies subsamples. We have also presented luminosity gap statistics, satellite fractions, and halo occupation numbers for the galaxies in our group sample. The main results are summarized as follows: \begin{enumerate} \item The conditional luminosity functions for central and satellite galaxies can be well modelled with a log-normal distribution and a modified Schechter form, respectively. The corresponding best fitting parameters are listed in Table 1. \item The average scatter of the log-normal luminosity distribution of central galaxies decreases from $\sim 0.15$ dex at the massive end ($\log [M_h/\msunh] \ga 13.5$) to $\sim 0.1$ dex at the low mass end ($\log [M_h/\msunh] \sim 12.0$). However, due to the method used to assign halo masses to the groups, at the low mass end this should be considered a lower limit on the true amount of scatter. \item The slope of the relation between the average luminosity of central galaxies (in the $^{0.1}r$-band) and halo mass, ${\rm d}\log L_c/{\rm d}\log M_h$, decreases from $\sim 0.68$ for $\log [M_h/\msunh] \ll 12.5$ to $\sim 0.17$ for$\log [M_h/\msunh] \gg 12.5$. For the stellar masses of the central galaxies we obtain that ${\rm d}\log M_{*,c}/{\rm d}\log M_h$, decreases from $\sim 1.83$ for $\log [M_h/\msunh] \ll 12.1$ to $\sim 0.22$ for$\log [M_h/\msunh] \gg 12.1$. \item The halo (group) occupation numbers of satellite galaxies accurately follow Poisson statistics. Since the same applies to dark matter sub-halos, this supports the standard picture that satellite galaxies are associated with dark matter sub-halos. \item In massive halos with masses $M_h\ga 10^{14}\msunh$ roughly 85\% (80\%) of the central (satellite) galaxies are red. These red fractions decrease to 50\% (40\%) in halos with masses $M_h \sim 10^{12}\msunh$. \item By comparing the scatter in the luminosities of BCGs to the luminosity difference between the BCG and its brightest satellite, we find that the BCGs form a `special' subclass, in that their luminosities can not be considered the extreme values of the distribution of satellite luminosities, expecially in halos with masses $M_h \la 10^{14.0} \msunh$. \item The fractions of fossil groups, which are defined as groups a with luminosity gap $\log L_1 - \log L_2 \ge 0.8$, decreases with increasing of halo mass from 18\%-60\% in halos with $M_h \sim 10^{13}\msunh$ to $\sim 2.5\%$ in halos with $M_h \sim 10^{14}\msunh$. \item The satellite fractions obtained from our group catalogue as functions of both luminosity and stellar mass (listed in Table~\ref{tab:fsat}) are in good agreement with independent constraints from analyses of galaxy clustering and galaxy-galaxy lensing. \end{enumerate} These results can be used to constrain the various physical processes related to galaxy formation and to interpret the various statistics used to describe large scale structures (e.g., galaxy correlation functions, pairwise velocity dispersions, etc.). Most of our findings are in good agreement with previous studies (e.g. Y05c, Zandivarez et al. 2006; Robotham et al. 2006) and can be linked to the semi-analytical modelling of galaxy formations (e.g., Kang et al. 2005; Zheng et al. 2005; Bower et al. 2006; Croton et al. 2006; De Lucia et al. 2007). The luminosity and stellar mass gap can be used to probe the specific formation properties of central galaxies (e.g., Vale \& Ostriker 2007). The fraction of the red and blue populations for central and satellite galaxies can be used to probe the color evolution of satellite galaxies (ven den Bosch et al. 2007b). | 7 | 10 | 0710.5096 |
0710 | 0710.3165_arXiv.txt | Stellar-mass black holes are discovered in X-ray emitting binary systems, where their mass can be determined from the dynamics of their companion stars\cite{rem06,cha06,oro03}. Models of stellar evolution have difficulty producing black holes in close binaries with masses \boldmath{$>10\,M_{\odot}$} (ref.\ 4), which is consistent with the fact that the most massive stellar black holes known so far\cite{cha06,oro03} all have masses within \boldmath{$1\sigma$} of \boldmath{$10\, M_{\odot}$}. Here we report a mass of \boldmath{$15.65 \pm 1.45\,M_{\odot}$} for the black hole in the recently discovered system M33 X-7, which is located in the nearby galaxy Messier 33 (M33) and is the only known black hole that is in an eclipsing binary\cite{pie06}. In order to produce such a massive black hole, the progenitor star must have retained much of its outer envelope until after helium fusion in the core was completed\cite{bro01}. On the other hand, in order for the black hole to be in its present 3.45 day orbit about its \boldmath{$70.0 \pm 6.9 M_{\odot}$} companion, there must have been a ``common envelope'' phase of evolution in which a significant amount of mass was lost from the system\cite{tau06}. We find the common envelope phase could not have occured in M33 X-7 unless the amount of mass lost from the progenitor during its evolution was an order of magnitude less than what is usually assumed in evolutionary models of massive stars\cite{sch92,mey94,vaz07}. | 7 | 10 | 0710.3165 |
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0710 | 0710.1998_arXiv.txt | The concept of black hole entropy is one of the most important enigmas of theoretical physics. It relates thermodynamics to gravity and allows substantial hints toward a quantum theory of gravitation. Although Bekenstein conjecture --assuming the black hole entropy to be a measure of the number of precollapse configurations-- has proved to be extremely fruitful, a direct and conclusive proof is still missing. This article computes accurately the entropy evaporated by black holes in $(4+n)$ dimensions taking into account the exact greybody factors. This is a key process to constrain and understand the entropy of black holes as the final state is unambiguously defined. Those results allow to generalize Zurek's important argument, in favor of the Bekenstein conjecture, to multi-dimensional scenarios. | 7 | 10 | 0710.1998 |
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0710 | 0710.3023_arXiv.txt | Three-dimensional large eddy simulations of solar surface convection using realistic model physics are conducted. The thermal structure of convective motions into the upper radiative layers of the photosphere, the range of convection cell sizes, and the penetration depths of convection are investigated. A portion of the solar photosphere and the upper layers of the convection zone, a region extending $60\times60$ Mm horizontally from 0 Mm down to 20 Mm below the visible surface, is considered. We start from a realistic initial model of the Sun with an equation of state and opacities of stellar matter. The equations of fully compressible radiation hydrodynamics with dynamical viscosity and gravity are solved. We use: 1) a high order conservative TVD scheme for the hydrodynamics, 2) the diffusion approximation for the radiative transfer, 3) dynamical viscosity from subgrid scale modeling. The simulations are conducted on a uniform horizontal grid of $600\times600$, with 168 nonuniformly spaced vertical grid points, on 144 processors with distributed memory multiprocessors on supercomputer MVS-15000BM in the Computational Centre of the Russian Academy of Sciences. | Convection near the solar surface has a strongly non-local and dynamical character. Hence, numerical simulations provide useful information on the spatial structures resulting from convection and help in constructing consistent models of the physical processes underlying the observed solar phenomena. We conduct an investigation of the temporal evolution and growth of convective modes on scales of mesogranulation and supergranulation in a three-dimensional computational box. In previous work by the author [\citet{ustyugs06}] it was shown that collective motion of small convective cells of granulation expels weak magnetic field on the edges of cells at mesogranular scales. The average size of such cells is 15-20 Mm and the lifetime of order 8-10 solar hours. Simulation of solar photosphere convection [\citet{stein06}] in a computational domain of size 48 Mm in the horizontal plane and 20 Mm in depth showed that the sizes of convective cells increase with depth. The purpose of this work is to investigate the development and scales of convection in a region of size 60 Mm in the horizontal plane and 20 Mm in depth. | 7 | 10 | 0710.3023 |
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0710 | 0710.3037_arXiv.txt | I present a new census of the stellar population in the Chamaeleon~I star-forming region. Using optical and near-IR photometry and followup spectroscopy, I have discovered 50 new members of Chamaeleon~I, expanding the census of known members to 226 objects. Fourteen of these new members have spectral types later than M6, which doubles the number of known members that are likely to be substellar. I have estimated extinctions, luminosities, and effective temperatures for the known members, used these data to construct an H-R diagram for the cluster, and inferred individual masses and ages with the theoretical evolutionary models of Baraffe and Chabrier. The distribution of isochronal ages indicates that star formation began 3-4 and 5-6~Myr ago in the southern and northern subclusters, respectively, and has continued to the present time at a declining rate. The IMF in Chamaeleon~I reaches a maximum at a mass of 0.1-0.15~$M_\odot$, and thus closely resembles the IMFs in IC~348 and the Orion Nebula Cluster. In logarithmic units where the Salpeter slope is 1.35, the IMF is roughly flat in the substellar regime and shows no indication of reaching a minimum down to a completeness limit of 0.01~$M_\odot$. The low-mass stars are more widely distributed than members at other masses in the northern subcluster, but this is not the case in the southern subcluster. Meanwhile, the brown dwarfs have the same spatial distribution as the stars out to a radius of $3\arcdeg$ (8.5~pc) from the center of Chamaeleon~I. | \label{sec:intro} The characteristics of the distributions of masses, ages, and positions in a newborn stellar population are determined by the process of star formation. As a result, measurements of these distributions in star-forming regions are potentially valuable for testing models of the birth of stars and brown dwarfs. For instance, the properties of the stellar initial mass function \citep[IMF,][]{mey00} can constrain the relative importance of turbulent fragmentation \citep{pn02}, gravitational fragmentation \citep{lar85}, dynamical interactions \citep{bon03}, and accretion and outflows \citep{af96} in regulating the final masses of stars. Models of the star formation rates of molecular clouds (e.g., constant, accelerating, bursts) can be tested against the distributions of ages and positions of members of young clusters \citep{fei96,pal97,har01}. The spatial distributions also provide insight into cloud fragmentation, binary formation, cluster dynamics, and the origin of brown dwarfs \citep{lar95,hil98,har02,luh06tau}. To obtain measurements of this kind, one must identify the members of star-forming regions and estimate their masses and ages. Only a few young stellar populations have been characterized in detail, such as the Orion Nebula Cluster \citep{hil97}, Taurus \citep{kh95}, and IC~348 in Perseus \citep{luh03ic}. The Chamaeleon~I star-forming region is amenable to a thorough census of its stellar population for several reasons. It is among the nearest star-forming regions \citep[$d=160$-170~pc,][]{whi97,wic98,ber99}, exhibits less extinction than many young clusters ($A_V\lesssim5$), is compact enough that a large fraction of the region can be surveyed to great depth in a reasonable amount of time, and is sufficiently rich for a statistically significant analysis of its stellar population. In addition, because Chamaeleon~I is relatively isolated from other star-forming regions, confusion with other young populations is minimal. Previous surveys have already identified more than 150 young stars and brown dwarfs in Chamaeleon~I through photometric variability, H$\alpha$ emission, X-ray emission, mid-infrared (IR) excess emission, and optical and near-IR color-magnitude diagrams \citep[][references therein]{com04,luh04cha,luh07cha}. However, in the census of known members produced by these surveys, the completeness as a function of mass and position is unknown \citep{luh04cha}. In this paper, I present a set of magnitude-limited surveys for members of Chamaeleon~I that have well-defined completeness limits (\S~\ref{sec:new}). I then use the new census of Chamaeleon~I to measure the star formation history (\S~\ref{sec:hr}), IMF (\S~\ref{sec:imf}), and spatial distribution of its stellar population (\S~\ref{sec:spatial}). | I have presented an extensive search for new members of the Chamaeleon~I star-forming region. Because the completeness limits of my survey are well-determined, I have been able to perform robust measurements of the distributions of members of Chamaeleon~I as a function of mass, position, and age. The primary results of this study are summarized as follows: \begin{enumerate} \item I have discovered 50 new members of Chamaeleon~I, which increases the census of known members to 226 objects. The new members include 14 objects that are later than M6 ($M\lesssim0.08$~$M_\odot$) and the two faintest known members of the cluster, which may have masses of only 0.005-0.01~$M_\odot$. The current census now contains 28 members that are likely to be substellar. \item The distribution of isochronal ages for members of Chamaeleon~I between 0.1-1~$M_\odot$ suggests that star formation has occurred for the past 3-4 and 5-6~Myr in the southern and northern subclusters, respectively, at rates that have declined with time. \item The IMF in Chamaeleon~I reaches a maximum at a mass of 0.1-0.15~$M_\odot$, which is similar to the turnover mass observed in IC~348 and the Orion Nebula Cluster \citep{hil97,hc00,mue02,mue03,luh03ic}. The substellar IMF is roughly flat in logarithmic units and shows no indication of reaching a minimum down to a completeness limit of 0.01~$M_\odot$. \item Chamaeleon~I does not contain a widely-distributed population of brown dwarfs, which is contrary to the predictions of some embryo ejection models. Instead, the substellar members share the same spatial distribution as the stars. However, low-mass stars in the northern subcluster do appear to have a wider distribution than members at other masses, which resembles the mass segregation that has been previously observed in Orion and IC~348 \citep{hc00,mue03}. \end{enumerate} | 7 | 10 | 0710.3037 |
0710 | 0710.1518_arXiv.txt | % After presenting three ways of defining a bulge component in disc galaxies, we introduce the various types of bulges, namely the classical bulges, the boxy/peanut bulges and the disc-like bulges. We then discuss three specific topics linked to bulge formation and evolution, namely the coupled time evolution of the bar, buckling and peanut strengths; the effect of velocity anisotropy on peanut formation; and bulge formation via bar destruction. | Three ways of defining a bulge have been used so far, one morphological, the second photometrical and the third kinematical. Based on morphology, a bulge is a component of a disc galaxy that has a smooth light distribution that swells out of the central part of a disc viewed edge-on. This definition has the disadvantage of being applicable only to edge-on systems and the advantage of necessitating only an image of the galaxy. The second definition is based on photometry and defines a bulge as the extra light in the central part of the galaxy, above the exponential profile fitting the remaining (non central) part of the disc. In earlier papers this component was fitted with an $r^{1/4}$ law, while more recent ones use its generalisation to an $r^{1/n}$ law (S\'ersic 1968). This definition has the advantage of being applicable to disc galaxies independent of their inclination. It has also the advantage of leading to quantitative results about the light distribution, but has the disadvantage of assigning to the bulge any extra central luminosity of the disc, independent of its origin. The third definition is based on kinematics, and in particular on the value of $V/\sigma$, or more specifically on the location of the object on the ($V/\sigma$, ellipticity) plot, which is often referred to as the Binney diagram (Binney 1978, 2005). This definition, potentially quite powerful, has unfortunately been very little used so far, due to the small number of galaxies for which the necessary data are available, a situation which is rapidly improving with large surveys, such as SAURON (Bacon \tal 2001; de Zeeuw \tal 2002; Peletier this volume). The lack of a single, clear-cut definition of a bulge is due to the fact that disc galaxies are viewed in different orientations and also to the fact that not all types of data are available for all objects. Nevertheless, it has led to considerable confusion and to the fact that bulges are an inhomogeneous class of objects. Indeed, many different types of objects, with very different properties and formation histories are included in the general term `bulges'. To remedy this, Kormendy (1993) and Kormendy \& Kennicutt (2004) distinguish classical bulges from pseudo-bulges, the latter category encompassing all bulges that are not classical. Athanassoula (2005a) argues that pseudo-bulges also are an inhomogeneous class of objects, and thus distinguishes three types of bulges. {\bf Classical bulges} are formed by gravitational collapse or hierarchical merging of smaller objects and corresponding dissipative gas processes. The material forming this bulge could be externally accreted, or could come from clumps in the proto-disc. In general, bulges of this type are formed before the actual disc (e.g. Steinmetz \& M\"uller 1995; Noguchi 1998; Immeli \tal 2004). Nevertheless, a bulge can also form from material externally accreted at much later stages (e.g. Pfenniger 1991; Athanassoula 1999; Aguerri, Balcells \& Peletier 2001; Fu, Huang, Deng 2003). Their morphological, photometrical and kinematical properties are similar to those of ellipticals. {\bf Box/peanut bulges} (B/P) form from a vertical instability of the disc material. This has often been observed in $N$-body simulations of bar-unstable discs, where the initial stage of bar formation is followed by a puffing up of the inner parts of the bar (e.g. Combes \& Sanders 1981; Combes \tal 1990; Raha \tal 1991; Athanassoula \& Misiriotis 2002; Athanassoula 2003, 2005a; O'Neil \& Dubinski 2003; Martinez-Valpuesta \& Shlosman 2004; Debattista \tal 2004, 2006; Martinez-Valpuesta \tal 2006). Viewed side-on (i.e. edge-on with the line of sight along the bar minor axis), this structure protrudes from the disc and has a characteristic boxy or peanut shape whose size is of the order of a few disc scale-lengths. Thus, a box/peanut bulge is just {\it part} of a bar seen side-on. Finally {\bf disc-like bulges} form from inflow of (mainly) gas material to the centre of the galaxy and subsequent star formation (e.g. Athanassoula 1992; Friedli \& Benz 1993; Heller \& Shlosman 1994; Wada \& Habe 1995) . The torque exerted by the bar pushes gas, and to a lesser extent also stars, to the inner parts of the disc where they form an inner disc. Star formation can be very high there, due to the increased gas density. Thus the result of this process should be a central disc, or disc-like object, whose stellar component should include a sizable fraction of young stars and whose size should be less than, or of the order of a kpc. Such a component could harbour sub-structures such as spirals, or bars, as is indeed sometimes observed (Kormendy \& Kennicutt 2004 and references therein). It is thus clear that disc-like bulges are very different objects from boxy/peanut bulges, since they are much smaller, have a different shape, different kinematics and provide a different type of excess on the radial photometric profiles. They also have different formation histories. The different formation histories of these three types of objects lead to different properties -- morphological, photometrical and kinematical -- which in turn help classify observed bulges into one of the three above mentioned types. Nevertheless, as stressed by Athanassoula (2005a), different types of bulges often co-exist and it is possible to find all three types of bulges in the same simulation, or galaxy. Realising the non-homogeneity of the objects classified as bulges and attempting to classify them is only the first step. Much more work is now necessary, particularly on two issues. The first one is the understanding of the formation and evolution of these types of objects. The second one is to predict the properties of these objects, starting from their formation scenarios. The latter is particularly important in order to bridge the gap between classification schemes based on formation histories and classification schemes based on observed properties. Here we make small contributions to both these issues, using $N$-body simulations. In Sect. 2 we present the time evolution of the bar, the buckling and the peanut strengths and their interplay. Sect. 3 discusses the velocity anisotropy and its link to the above strengths. Finally, in Sect. 4 we discuss the photometrical properties of a destroyed bar and boxy/peanut bulge. | 7 | 10 | 0710.1518 |
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0710 | 0710.3347_arXiv.txt | It has been suggested that fossil groups could be the canibalized remains of compact groups, that lost energy through tidal friction. However, in the nearby universe, compact groups which are close to the merging phase and display a wealth of interacting features (such as HCG 31 and HCG 79) have very low velocity dispersions and poor neighborhoods, unlike the massive, cluster-like fossil groups studied to date. In fact, known z$=$0 compact groups are very seldom embedded in massive enough structures which may have resembled the intergalactic medium of fossil groups. In this paper we study the dynamical properties of CG6, a massive compact group at z$=$0.220 that has several properties in common with known fossil groups. We report on new \sloang~and \sloani~imaging and multi-slit spectroscopic performed with GMOS on Gemini South. The system has 20 members, within a radius of 1 h$_{70}^{-1}$ Mpc, a velocity dispersion of 700~\kms~and has a mass of 1.8 $\times$ 10$^{14}$ h$_{70}^{-1}$~\Msol, similar to that of the most massive fossil groups known. The merging of the four central galaxies in this group would form a galaxy with magnitude $M_{r'} \sim -23.4$, typical for first-ranked galaxies of fossil groups. Although nearby compact groups with similar properties to CG 6 are rare, we speculate that such systems occurred more frequently in the past and they may have been the precursors of fossil groups. | \label{sec:intro} Groups of galaxies are small systems of typically a few L$^{*}$ galaxies, which comprise over 55\% of the nearby structures in the universe. A small fraction of galaxy groups are classified as {\it compact} groups, which are responsible for $\sim$ 1\% of the luminosity density of the universe \citep{mdoh91}. Although they are rare objects in the nearby universe, their high galactic densities and low velocity dispersions make them ideal systems for the study of galaxy transformation through galaxy-galaxy collisions. As expected, these systems have a high fraction of interacting members, although merged objects are rare \citep{zepf93}. They are commonly believed to evolve through dynamical friction and finally merge to form one single galaxy \citep{Barnes92}. \citet{vikhlinin99} and \citet{jones03} have suggested that the merging of compact groups can lead to the formation of fossil groups. A fossil group (FG, hereafter) is a system with an extended and luminous X-ray halo (L$_X$ $>$ $10^{42}$ h$_{70}^{-2}$ erg s$^{-1}$), dominated by one single brighter than L$^{*}$ elliptical galaxy, surrounded by low-luminosity companions \citep[where the difference in magnitude between the bright dominant elliptical and the next brightest companion is $>$ 2 mag in the R-band;][]{jones03}. One important goal of this article is to investigate if compact groups (CG, hereafter) as we known them today, could be the precursors of FGs. In order to answer this question, we summarize, in section 2, the properties of a few of the most strongly interacting nearby CGs known, which are about to merge, and in section 3, we describe the properties of the five FGs which have been studied spectroscopically so far. In section 4, we present new observations for a CG embedded in a cluster-size potential, at redshift z$=$0.22 and section 5 puts together all the observations described and discusses the CG-FG scenario. Throughout this paper we adopt when necessary a standard cosmological model: $H_{0}=70\,h_{70}\,$km~s$^{-1}$~Mpc$^{-1}$, $\Omega_{m}=0.3$ and $\Omega_{\Lambda}=0.7$. At z$=$0.22, 1\arcsec~corresponds to $ 3.5\, h_{70}^{-1}\,$kpc. \section {Interacting compact groups: HCG 31, HCG 79 and HCG 92} \label{sec:hcg} There is evidence from both observations and simulations that groups evolve through dynamical friction and coalesce to form more compact structures as the universe ages. A few of the most compact, and therefore most evolved groups known, from Hickson's catalogue \citep{h92} are HCG 31, HCG 79 (or Seyfert Sextet) and HCG 92 (or Stephan's quintet). The study of these groups is important to help understanding processes common in merging systems, environments that may have occurred more often in the high-redshift universe. HCG 31 is a group at z$\sim$0.013 and with a velocity dispersion of $\sigma$ $\sim$ 60 \kms. This is a gas-rich group with intense star forming activity \cite[e.g. ][]{moetal06, amram07}, dominated by a central pair of interacting dwarf galaxies A$+$C. HCG 31 is thought to be in a pre-merging phase \citep{amram05,vm05} and it has well developed tidal tails seen in H$\alpha$ and HI. The group hosts two excellent candidates for tidal dwarf galaxies, namely member F, in the southern tail and member R, 50 h$_{70}^{-1}$ kpc to the north of the group (for an assumed distance modulus of DM=33.8). HCG 79, also known as {\it Seyfert Sextet}, was originally identified as a sextet of galaxies but it is now known to be a quartet at z$=$ 0.0145 (the 5th object is in the background and the 6th is a luminous tidal debris to the northeast of the group). This is the most CG in Hickson's catalogue with a galaxy-galaxy distance below 10 kpc (for an adopted DM=34.0) and a velocity dispersion of $\sigma$ = 138 \kms. The four galaxies present morphological distortions and increased activity (tidal debris, bar in HCG 79B, dust lane in HCG 79A, radio and infrared emission, disturbed rotation curves and nuclear activity). The group presents a prominent intra-group light envelope which contains 45\% of the total light of the group \citep{darocha05} and irregular envelopes of HI \citep{wil91} and X--rays \citep{pil95}. These suggest that recent or on-going interaction is taking place within this system. HCG 92, also known as Stephan's quintet, is in reality a quartet with z$=$ 0.0215 and a foreground galaxy. It is the most well studied CG -- multi wavelength data are available from radio to X-rays. Most of the gaseous material in Stephan's quintet is concentrated not in the galaxies but in the intragroup medium, suggesting that collisions among group members may have happened frequently. A number of tidal dwarf galaxies and intergalactic HII regions have been identified in this group \citep[e.g.][]{moetal01,moetal04,xu05}. Of the three groups described above, HCG 92 is the only one to have detected X-rays, with a total bolometric luminosity of 2.96$\times$10$^{42}$ h$^{-2}_{70}$ \ergs~\citep{Xue2000}. These three spiral-rich groups are thought to be in their final stages of evolution -- they are, in fact, some of the most compact systems found in the Hickson's catalogue. Yet, they have members that can be clearly identified as individual galaxies, suggesting that once merging starts, it may proceed quickly, and the groups may no longer be recognized as such. The bright members of these groups will almost certainly end up as a single galaxy pile. A discussion of whether these systems will most likely end up as FGs or as single isolated elliptical galaxies is deferred to section 5. In the following section, some of the optical properties of the FGs studied so far are summarized. | Dynamical friction and subsequent merging are probably the processes responsible for the lack of L$^{*}$ galaxies in FGs. Considering the merging scenario, it is possible that the overluminous central galaxy in a FG has been formed within a substructure, inside a larger structure. In that case, one could think of a scenario where a CG was formed within a rich group, which would then have merged, leaving behind the brightest elliptical galaxy of what today is seen as a FG. One weak argument against this scenario is that the nearby examples of CGs are not usually found within such massive structures, but instead are more often surrounded by very sparse structures. There are, however, examples such as CG6, surrounded by large numbers of lower-luminosity galaxies, inhabiting a deep potential well. We would like to test the hypothesis that CGs, as observed in the nearby universe, could be the precursors of FGs. We may examine two aspects: (1) if the sum of the luminosities of the brightest CG galaxies is similar to the luminosity of a first-ranked FG galaxy and; (2) if the neighbourhoods of CGs are rich, i.e., if the system as a whole (group plus environment) has a velocity dispersion/mass similar to that of a FG. We compute the total luminosity of the galaxies in the soon-to-merge CGs, HCG 31, HCG 79 and HCG 92, to check how these compare with the luminosities of first-ranked galaxies in FGs. Adding up the luminosities of galaxies HCG 31 A to C, G and Q, which are the brightest in the group HCG 31, a magnitude of M$_R = -22.5$ is obtained (for a distance modulus, DM = 33.8). Summing up the luminosities of galaxies HCG 79 A-D, an equal total magnitude of M$_R=-22.5$ is obtained (for DM = 34.0). These are upper limits on the luminosities of these objects given that several members are starburst galaxies. After fading, the merged central object in HCG 31 and HCG 79 will have somewhat lower magnitudes than that of a typical first-ranked galaxy in a FG. Fossil groups first-ranked galaxies have luminosities well above L$^{*}$. For the five FGs studied by \citet{jones03}, the first-ranked galaxies had a median luminosity of $M_R=-23.2$ and for the 34 FGs found in the SDSS DR5 by \citet{dosSantos07} the median luminosity was $M_R=-23.5$. Although for HCG 92, the final object (adding up luminosities of galaxies A-E) would have an absolute magnitude of $M_R=-24.2$ (for DM$=$34.8), which after allowing for some fading, could be similar to that of an FG first-ranked galaxy, HCG 92 would possibly still not resemble an FG when merged, because its neighbourhood is very sparse, i.e., it is not embedded in any larger structure, as it is often the case for the central galaxy in FGs. This is in agreement with its relatively low bolometric X-ray luminosity of 2.96 $\times$ 10$^{42}$ h$_{70}^{-2}$ ergs s$^{-1}$ \citep{Xue2000}. The environments of nearby CGs have been surveyed by \citet{ribeiro98,zabludoff98,carrasco06} among others. Spectroscopy of dozens of members in the neighbourhood of quite a number of groups was obtained, confirming in all cases that CGs have low velocity dispersions typical of the group regime (typically 200-300 km s$^{-1}$). In fact, even for HCG 62, thought to be one of the most massive CGs in Hickson's catalogue, the velocity dispersion obtained from 45 members of the system showed that it is a bonafide group (376 km s$^{-1}$). HCG 62 was suggested by \citet{ponman93} as an example of a system that could turn into a FG in a few Gygayears, but its velocity dispersion is still much lower than the value of $\sim$ 600 km s$^{-1}$, typical for rich FGs. Two other massive nearby CGs in Hickson's catalogue are HCG 94 and HCG 65. The first is known to have a very high bolometric X-ray luminosity of 2.35 $\times$ 10$^{44}$ h$_{70}^{-2}$ \citep{Xue2000} which may, however, be contaminated by the emission of a nearby cluster. The velocity dispersion obtained from 11 members in this system gives a value of 479 km s$^{-1}$. HCG 65 is the center of the cluster Abell 3559. It is in the heart of the Shapley supercluster and its location makes it hard to disentangle its dynamics and determine its mass. The three most massive Hickson CGs known, HCG 62, HCG 65 and HCG 94, are strongly early-type dominated, as expected from the velocity dispersion-morphology relation observed for CGs. The conclusion is then that a typical CG, as observed at z=0, is unlike to turn into a FG. It is more likely to merge into an isolated elliptical galaxy. For the compact group CG 6, at $z=0.22$ and $\sigma =$703 km s$^{-1}$, if we merge the four central galaxies (A, B, C and D), we end up with a galaxy with total magnitude \sloani$=$16.31 and \sloang$=$17.80 mag (M$_{i'}$=--23.87 and M$_{g'}$=--22.38, with no k-correction). Using the color relation for a galaxy at z$=$0.2 from \cite{fuk95}, the magnitude in r$^{'}$ will be 16.83 or M$_{r}=-23.35$. This magnitude is similar to those of typical central galaxies in FGs and the velocity dispersion of the system is typical for the studied FGs. However, no gap of at least two magnitudes between the first-ranked relic and the remaining objects of the system would be observed because there is at least one other bright galaxy in the system within half the virial radius of the group. We point out that one other example of a possible massive system, at z=0.39, which may turn into a FG, has recently been discovered by \cite{rines07}. Spectroscopic studies of CGs at medium redshifts may find many more of such objects. | 7 | 10 | 0710.3347 |
0710 | 0710.1497_arXiv.txt | We present the results of an analysis of a well-selected sample of galaxies with active and inactive galactic nuclei from the Sloan Digital Sky Survey, in the range $0.01 < z < 0.16$. The SDSS galaxy catalogue was split into two classes of active galaxies, Type~2 AGN and composites, and one set of inactive, star-forming/passive galaxies. For each active galaxy, two inactive control galaxies were selected by matching redshift, absolute magnitude, inclination, and radius. The sample of inactive galaxies naturally divides into a red and a blue sequence, while the vast majority of AGN hosts occur along the red sequence. In terms of H$\alpha$ equivalent width, the population of composite galaxies peaks in the valley between the two modes, suggesting a transition population. However, this effect is not observed in other properties such as colour-magnitude space, or colour-concentration plane. Active galaxies are seen to be generally bulge-dominated systems, but with enhanced H$\alpha$ emission compared to inactive red-sequence galaxies. AGN and composites also occur in less dense environments than inactive red-sequence galaxies, implying that the fuelling of AGN is more restricted in high-density environments. These results are therefore inconsistent with theories in which AGN host galaxies are a `transition' population. We also introduce a systematic 3D spectroscopic imaging survey, to quantify and compare the gaseous and stellar kinematics of a well-selected, distance-limited sample of up to 20 nearby Seyfert galaxies, and 20 inactive control galaxies with well-matched optical properties. The survey aims to search for dynamical triggers of nuclear activity and address outstanding controversies in optical/IR imaging surveys. | Active Galactic Nuclei (AGN) have long been considered a curiosity in their own right \citep{schmidt63, lybell69,blan74}, but are now recognised to be integral to galaxy formation and evolution. At early cosmological epochs, gas-rich galaxies are thought to form and collide in violent mergers \citep{mihos96}, triggering vast bursts of star formation, and fuelling supermassive black holes in their cores \citep{kriek06}. Recent simulations of isolated and merging galaxies by \cite{spring05} have incorporated feedback from star formation and black-hole accretion, and find that once an accreting supermassive black hole (SMBH) has grown to some critical size, the AGN feedback terminates its growth as a large fraction of the remaining nuclear gas is driven out by the powerful quasar. In the current epoch, the peak of the quasar era is over and the galaxy merger rate has declined \citep{struck97}. Nevertheless, at least 20\% of today's galaxies show scaled-down quasar activity in their centres, and direct measurements of active and inactive galaxy dynamics have revealed a tight correlation between the central black-hole mass and host galaxy stellar velocity dispersion or bulge mass. This $M_{\bullet}-\sigma$ relation \citep{geb00,merr01} points to an intimate link between host galaxy evolution and central black-hole growth, suggesting that all bulge-dominated galaxies today harbour dead quasars and that a lack of nuclear activity cannot be attributed to the absence of a central black hole. Therefore, given the ubiquity of supermassive black holes, what determines the degree of nuclear activity in todays galaxies and what role is played by the host galaxy in triggering and fuelling their dormant black holes remain open issues. Previously, studies of nearby active galaxies were based on small galaxy samples, like the C$f$A Seyfert Sample \citep{huchra92}, the 12 Micron Active Galaxy Sample \citep{rush93} and the \cite{ho95} galaxy sample of approximately 486 galaxies covering a range of activity types. Now, however, the standard has been set by the Sloan Digital Sky Survey (SDSS), from which a range of useful samples of both active and inactive galaxies can be selected in a well-defined, uniform way. The SDSS \citep{york00} provides, for the first time, 5-band photometry and spectroscopy of many thousands of low redshift AGN \citep{kauf03c,hao05a}, enabling constraints to be placed on galaxy and AGN evolution over a wide range of galaxy masses, and acts as the definitive supporting data to all detailed follow-up galaxy studies. Initial investigations into the growth and evolution of black holes using these databases have yielded contrasting results --- the origin of the correlation between galaxy bulge and central black-hole masses is hotly debated in particular. \citet{heck04} find that the majority of present-day accretion occurs onto 10$^{8}$ solar-mass black holes in moderate mass galaxies, suggesting that bulge and black-hole evolution is still tightly coupled today, and that the evolution of AGN luminosity functions is driven by a decrease in the mass scale of accreting black holes. In contrast, the \citet{hao05b} study of about 3\,000 SDSS AGN concludes that evolution in AGN luminosity functions is driven by evolution in the Eddington ratio, rather than black-hole mass. Meanwhile, \cite{grup04} and \cite{grup05} use a sample of 75 X-ray selected AGN to argue that black holes in Narrow Line Seyfert 1 galaxies in particular, grow by accretion in well-formed bulges to produce the $M_{\bullet}-\sigma$ relation over time, refuting theories for the origin of black-hole / bulge relations in which black-hole mass is a constant fraction of bulge mass at all epochs or in which bulge growth is controlled by AGN feedback \citep{king03,ferr02,kriek06}. The subclass of NLS1s has been further explored more recently by \cite{komo07}, who find that NLS1s are accreting at a rate higher than the Eddington rate, confirming that their BHs must be growing. They suggest that either NLS1 galaxies evolve into Broad Line Seyfert 1 (BLS1) galaxies with respect to their black hole mass distribution, which would require some change in the bulge properties, possibly due to feedback, or that NLS1s are just low-mass extensions of BLS1 galaxies, and the high accretion rate could just be caused by a relatively short-lived accretion phase. Another characteristic revealed by large automated surveys such as the SDSS, is that the galaxy population is found to be bimodal in colour \citep{lilly95,strat01}, with it now being more natural to describe a galaxy as being on the ``red sequence'' or ``blue sequence'', rather than being ``early type'' or ``late type'' \citep{bald06}. A key goal of galaxy evolution theory is then to explain the colour bimodality of galaxies, the relationships within each sequence, and where active hosts fit into this picture. Therefore, the understanding of galaxy formation and evolution processes necessitates full inclusion of AGN and their hosts. Recently efforts have been made to try to understand where AGN host galaxies fall in colour-magnitude space. The red sequence consists of mainly massive, passively evolving galaxies, while the majority of galaxies show blue colours, attributed to ongoing star formation. The two sequences are separated by a relatively narrow valley in colour space. The emerging consensus is that intense star formation is fuelled by galaxy mergers at high redshift, which forms massive bulges, and at some point the star formation ceases, resulting in the galaxy migrating from the blue sequence to the red sequence. AGN have been singled out as the mechanism for star-formation quenching, leading to suggestions that AGN should occupy a distinct, `transition' region, of the colour-magnitude diagram (CMD). However this does not necessarily have to be the case, as the resulting galaxy could just be bluer due to increased star-burst activity, or redder due to enhanced dust (e.g., Luminous Infra-Red Galaxies), or some combination of the original colours of the merging galaxies. Young stellar populations and nuclear starbursts are therefore also important components in the dynamics of nuclear activty \citep{gonza98,sarzi07}. \cite{nand07} studied the colour-magnitude relation for a sample of 50 X-ray selected AGN from the AEGIS survey (All-wavelength Extended Groth strip International Survey), in the range $0.6 < z < 1.4$. They conclude that AGN fall on the red sequence or on the red edge of the blue sequence, with many in between these two modes. \cite{mart07} further explored the idea of a `transition' region by exploiting the separation of the blue and red sequences in the ${\rm UV} - r$ colour-magnitude diagram. Using a sample of UV selected galaxies from the GALEX Medium Imaging Survey, along with SDSS data, they explored the nature of galaxies in the transition zone. The AGN fraction of their sample was found to peak in the transition zone, and they also found circumstantial evidence that star-formation quenching rates were higher in higher luminosity AGN. Higher image quality than the SDSS is provided by the Millennium Galaxy Catalogue (MGC; \citealt{liske03,cross04}), which has obtained redshifts for 10\,095 galaxies to $B < 20 {\rm\,mag}$, covering 37 deg$^{2}$ of equatorial sky. The MGC study of galaxy bimodality was in the colour-structure plane, (\citealt{drive06}; hereafter D06) and expressed a contrasting theory for the bimodality. The MGC survey has sufficient resolution to achieve reliable bulge-disc decomposition \citep{allen06}, and in separating out the two components, D06 claim that galaxy bimodality is caused by the bulge-disc nature of galaxies, and not by two distinct galaxy classes at different evolutionary stages. In this case, they show that the bulge-dominated, early-type galaxies populate one peak and the bulge-less, late-type galaxies occupy the second. The early- and mid-type spirals sprawl across and between the peaks. They also propose that the reason for the dual structure of galaxies is due to galaxy formation proceeding in two stages; first, there is an initial collapse phase, which forms the centrally concentrated core and a black hole, and second, there is the formation of a planar rotating disc caused by accretion of external material building up the galaxy disk. In this paper, we present the properties of a magnitude-limited sample of ``active'' and ``inactive'' galaxies carefully selected from the SDSS, which also acts as the parent sample for detailed followup studies using the IMACS-IFU on the Magellan 6.5m telescope. Our sample selection is described in \S 3, the sample properties in \S 4, and implications of these properties in \S 5. We conclude by introducing the Magellan survey, which is now underway. | We have carried out a robust classification of `active' and `inactive' galaxies in the SDSS, based on their emission line properties and locations in the BPT diagram. We have compared the active galaxies with well selected control galaxies matched in redshift, absolute magnitude, aspect ratio and radius. We find that: \begin{itemize} \item Type~2 AGN host galaxies occur mainly on the red sequence of the CMD, but show increased levels of H$\alpha$ flux in their emission compared to inactive red-sequence galaxies (Fig.~\ref{fig:H-alpha}(a) and Fig.~\ref{fig:CMD}(a)); \item A separate class of composite galaxies appears to peak on the blue edge of the red sequence on the CMD, whereas the peak of the H$\alpha$ distribution places composite galaxies firmly in the valley between the blue and red sequences (Fig.~\ref{fig:H-alpha}(b) and Fig.~\ref{fig:CMD}(b)); \item Colour-concentration relations, however, show a more complex, possibly double, morphology in the peak of the composite distribution, rather than being in a valley (Fig.~\ref{fig:colour-conc-comp}); \item AGN (and composites) are found in less dense environments on average then matched inactive red-sequence galaxies. The more clustered inactive galaxies are likely to be satellite galaxies in high-density environments (Fig.~\ref{fig:mean-enviro} and Fig.~\ref{fig:distrib-enviro}). \end{itemize} The key to understanding this in more detail now lies in the dynamics of the central regions of galaxies, to understand what activates some galaxies, but not others. | 7 | 10 | 0710.1497 |
0710 | 0710.4567_arXiv.txt | We construct merger trees from the largest database of dark matter haloes to date provided by the Millennium simulation to quantify the merger rates of haloes over a broad range of descendant halo mass ($10^{12} \la M_0 \la 10^{15} M_\odot$), progenitor mass ratio ($10^{-3} \la \xi \le 1$), and redshift ($0 \le z \la 6$). We find the mean merger rate {\it per halo}, $B/n$, to have very simple dependence on $M_0$, $\xi$, and $z$, and propose a universal fitting form for $B/n$ that is accurate to 10-20\%. Overall, $B/n$ depends very weakly on the halo mass ($\propto M_0^{0.08}$) and scales as a power law in the progenitor mass ratio ($\propto \xi^{-2}$) for minor mergers ($\xi \la 0.1$) with a mild upturn for major mergers. As a function of time, we find the merger rate per Gyr to evolve roughly as $(1+z)^{n_m}$ with $n_m=2-2.3$, while the rate per unit redshift is nearly independent of $z$. Several tests are performed to assess how our merger rates are affected by, e.g. the time interval between Millennium outputs, binary vs multiple progenitor mergers, and mass conservation and diffuse accretion during mergers. In particular, we find halo fragmentations to be a general issue in merger tree construction from $N$-body simulations and compare two methods for handling these events. We compare our results with predictions of two analytical models for halo mergers based on the Extended Press-Schechter (EPS) model and the coagulation theory. We find that the EPS model overpredicts the major merger rates and underpredicts the minor merger rates by up to a factor of a few. | In hierarchical cosmological models such as $\Lambda$CDM, galaxies' host dark matter haloes grow in mass and size primarily through mergers with other haloes. As the haloes merge, their more centrally concentrated baryonic components sink through dynamical friction and merge subsequently. The growth of stellar masses depends on both the amount of mass brought in by mergers and the star formation rates. Having an accurate description of the mergers of dark matter haloes is therefore a key first step in quantifying the mergers of galaxies and in understanding galaxy formation and growth. Earlier theoretical studies of galaxy formation typically relied on merger trees generated from Monte Carlo realisations of the merger rates given by the analytical extended Press-Schechter (EPS; \citealt{1993MNRAS.262..627L, 1991ApJ...379..440B}) model (e.g. \citealt{1993MNRAS.264..201K, 1999MNRAS.310.1087S, 2000MNRAS.319..168C}). Some recent studies have chosen to bypass the uncertainties and inconsistencies in the EPS model by using halo merger trees from $N$-body simulations directly (\citealt{1999MNRAS.303..188K, 2000MNRAS.311..793B, 2003MNRAS.338..903H, 2005ApJ...631...21K, 2005Natur.435..629S}). As we find in this paper, obtaining robust halo merger rates and merger trees requires rich halo statistics from very large cosmological simulations as well as careful treatments of systematic effects due to different algorithms used for, e.g., assigning halo masses, constructing merger trees, removing halo fragmentation events, and choosing time spacings between simulation outputs. The aim of this paper is to determine the merger rates of dark matter haloes as a function of halo mass, merger mass ratio (i.e. minor vs major), and redshift, using numerical simulations of the $\Lambda$CDM cosmology. This basic quantity has not been thoroughly investigated until now mainly because large catalogues of haloes from finely spaced simulation outputs are required to provide sufficient merger event statistics for a reliable construction of merger trees over a wide dynamic range in time and mass. We achieve this goal by using the public database of the Millennium simulation \citep{2005Natur.435..629S}, which follows the evolution of roughly $2\times10^{7}$ dark matter haloes from redshift $z=127$ to $z=0$. This dataset allows us to determine the merger rates of dark matter haloes ranging from galaxy-mass scales of $\sim 10^{12} M_\odot$ over redshifts $z=0$ to $\sim 6$, to cluster-mass scales up to $\sim 10^{15} M_\odot$ for $z=0$ to a few. We are also able to quantify the merger rates as a function of the progenitor mass ratio $\xi$, from major mergers ($\xi \ga 0.1$) down to minor mergers of $\xi \sim 0.03$ for galaxy haloes and down to $\xi \sim 3\times 10^{-4}$ for cluster haloes. The inputs needed for measuring merger rates in simulations include a catalogue of dark matter haloes and their masses at each redshift, and detailed information about their ancestry across redshifts, that is, the merger tree. Unfortunately there is not a unique way to identify haloes, assign halo masses, and construct merger trees. In this paper we primarily consider a halo mass definition based on the standard friends-of-friends (FOF) algorithm and briefly compare it with an alternative mass definition based on spherical overdensity. For the merger trees, we investigate two possible algorithms for treating events in which the particles in a given progenitor halo end up in more than one descendant halo ('fragmentations'). We find that these events are common enough that a careful treatment is needed. In the conventional algorithm used in the literature, the progenitor halo is linked one-to-one to the descendant halo that has inherited the largest number of the progenitor's particles. The ancestry links to the other descendant haloes are severed (for this reason we call this scheme 'snipping'). We consider an alternative algorithm ('stitching') in this paper, in which fragmentations are assumed to be artefacts of the FOF halo identification scheme. We therefore choose to recombine the halo fragments and stitch them back into the original FOF halo. Earlier theoretical papers on merger rates either relied on a small sample of main haloes to estimate the overall redshift evolution over a limited range of halo masses, or were primarily concerned with the mergers of {\it galaxies} or {\it subhaloes}. For halo mergers, for example, \cite{1999AJ....117.1651G} studied $z < 1$ major mergers of galaxy-sized haloes in an open CDM and a tilted $\Omega_m=1$ CDM model using $N$-body simulations in a 100 Mpc box and $144^3$ particles. \cite{2001ApJ...546..223G} used a sample of $\sim 4000$ haloes to study the environmental dependence of the redshift evolution of the major merger rate at $z < 2$ in $\Lambda$CDM. \cite{2006ApJ...652...56B} studied major mergers of subhaloes in $N$-body simulations in a 171 Mpc box with $512^3$ particles and the connection to the observed close pair counts of galaxies. For galaxy merger rates, \cite{2002ApJ...571....1M} and \cite{2006ApJ...647..763M} are based on up to $\sim 500$ galaxies formed in SPH simulations in $\sim 50$ Mpc boxes with up to $144^3$ gas particles, while \cite{2007arXiv0708.1814G} used the semi-analytical galaxy catalogue of \cite{2006MNRAS.366..499D} based on the Millennium simulation. This paper is organised as follows. Section~\ref{MillenniumSection} describes the dark matter haloes in the Millennium simulation (\S\ref{DarkMatterHaloes}) and how we construct the merger trees (\S\ref{mergertrees}) . We then discuss the issue of halo fragmentation and the two methods ('snipping' and 'stitching') used to treat these events in \S\ref{FOFLims}. The notation used in this paper is summarised in \S\ref{Notation}. Section~\ref{StatsSection} describes how mergers are counted (\S\ref{MultiCount}) and presents four (related) statistical measures of the merger rate (\S\ref{stats}). The relation between these merger rate statistics and the analytical merger rate based on the Extended Press-Schechter (EPS) model is derived in Section~\ref{EPSConnection}. Our main results on the merger rates computed from the Millennium simulation are presented in Section~\ref{Results}. We first discuss the $z\approx 0$ results and quantify the merger rates as a function of the descendant halo mass and the progenitor mass ratios using merger trees constructed from the stitching method (\S\ref{MRz0}). The evolution of the merger rates with redshifts up to $z\sim 6$ is discussed in Section~\ref{redshiftdependence}. We find a simple universal form for the merger rates and present an analytic fitting form that provides a good approximation (at the 10-20\% level) over a wide range of parameters (\S\ref{FitSection}). Section~\ref{Tests} compares the stitching and snipping merger rates (\S\ref{SnipvsStitch}) and presents the key results from a number of tests that we have carried out to assess the robustness of our results. Among the tests are: time convergence and the dependence of the merger rates on the redshift spacing $\dz$ between the Millennium outputs used to construct the merger tree (\S\ref{TimeResolutionSection}); how the counting of binary vs multiple progenitor mergers affects the merger rates (\S\ref{MultiCountValid}); mass non-conservation arising from 'diffuse' accretion in the form of unresolved haloes during mergers (\S\ref{MassCons}); and how the definition of halo masses and the treatment of fragmentation events affect the resulting halo mass function (\S\ref{MassFunction}). In Section~\ref{DiscussionSection}, we discuss two theoretical frameworks that can be used to model halo mergers: EPS and coagulation. A direct comparison of our merger rates and the EPS predictions for the Millennium $\Lambda$CDM model shows significant differences over a large range of parameter space (\S\ref{eps}). Section~\ref{Coagulation} discusses Smoluchowski's coagulation equation and the connection between our merger rates and the coagulation merger kernel. The appendix compares a third merger tree (besides snipping and stitching) constructed from the Millennium catalogue by the Durham group \citep{2006MNRAS.370..645B, 2006MNRAS.367.1039H,2003MNRAS.338..903H}. Two additional criteria are imposed on the subhaloes in this algorithm to reduce spurious linkings of FOF haloes. We find these criteria to result in reductions in both the major merger rates and the halo mass function. The cosmology used throughout this paper is identical to that used in the Millennium simulation: a $\Lambda\textrm{CDM}$ model with $\Omega_m=0.25$, $\Omega_b=0.045$, $\Omega_\Lambda=0.75$, $h=0.73$, an initial power-law index $n=1$, and $\sigma_8=0.9$ \citep{2005Natur.435..629S}. Masses and lengths are quoted in units of $M_\odot$ and Mpc without the Hubble parameter $h$. | In this paper we have computed the merger rates of dark matter FOF haloes as a function of descendant halo mass $M_0$, progenitor mass ratio $\xi$, and redshift $z$ using the merger trees that we constructed from the halo catalogue of the Millennium simulation. Our main results are presented in Figs.~\ref{fig:B} to \ref{fig:Rz}, which show very simple and nearly separable dependence on $M_0$, $\xi$, and $z$. The mean merger rate per descendant FOF halo, $B/n$, is seen to depend very weakly on the halo mass $M_0$ (Fig.~\ref{fig:B} right panel and Fig.~\ref{fig:Rmass}). As a function of redshift $z$, the per halo merger rate in units of per Gyr increases as $(1+z)^\alpha$, where $\alpha\sim 2$ to 2.3 (top panel of Fig.~\ref{fig:Rz}), but when expressed in units of per redshift, the merger rate depends very weakly on $z$ (bottom panel of Fig.~\ref{fig:Rz}). Regardless of $M_0$ and $z$, the dependence of $B/n$ on the progenitor mass ratio, $\xi = M_i/M_1$, is a power law to a good approximation in the minor merger regime ($\xi \la 0.1$) and shows an upturn in the major merger regime (Fig.~\ref{fig:B}). These simple behaviours have allowed us to propose a universal fitting formula in equation~(\ref{fiteqn}) that is valid for $10^{12}\leq M_0\la 10^{15} M_\odot$, $\xi \ga 10^{-3}$, and up to $z\sim 6$. Throughout the paper we have emphasised and quantified the effects on the merger rates due to events in which a progenitor halo fragments into multiple descendant haloes. We have shown that the method commonly used to remove these fragmented haloes in merger trees -- the snipping method -- has relatively poor $\dz$-convergence (Figs.~\ref{fig:TR} and \ref{fig:RzTR}). Our alternative approach -- the stitching method -- performs well with regards to this issue without drastically modifying the mass conservation properties or the mass function of the Millennium FOF catalogue (Figs.~ \ref{fig:DM} and \ref{fig:MassFunction}). We have computed the two predictions for merger rates from the analytical EPS model for the same $\Lambda$CDM model used in the Millennium simulation. At $z=0$, we find the EPS major merger rates to be too high by 50-100\% (depending on halo mass) and the minor merger rates to be too low by up to a factor of 2-5 (Fig.~\ref{fig:EPS}). The discrepancy increases at higher $z$. The coagulation equation offers an alternative theoretical framework for modelling the mergers of dark matter haloes. We have discussed how our merger rate is related to the coagulation merger kernel in theory. In practice, however, we find that mergers in simulations are not always mass-conserving binary events, as assumed in the standard coagulation form given by equation~(\ref{eqn:coag}). Equation~(\ref{eqn:coag}) will therefore have to be modified before it can be used to model mergers in simulations. \cite{2001ApJ...546..223G} studied the rate of major mergers (defined to be $\xi \ge 1/3$ in our notation) in $N$-body simulations and found a steeper power law dependence of $\propto (1+z)^3$ (at $z\la 2$) for the merger rate per Gyr than ours. Their simulations did not have sufficient mass resolution to determine the rate at $z \ga 2$. It is important to note, however, that our $B/n$ at redshift $z$ measures the instantaneous rate of mergers during a small $\Delta z$ interval at that redshift. By contrast, they studied the merging history of {\it present-day} haloes and measured only the rate of major mergers for the most massive progenitor at redshift $z$ of a $z=0$ halo (see their paragraph 4, section 2). A detailed comparison is outside the interest of this paper. Mergers of dark matter haloes are related to but not identical to mergers of galaxies. It typically takes the stellar component of an infalling galaxy extra time to merge with a central galaxy in a group or cluster after their respective dark matter haloes have been tagged as merged by the FOF algorithm. This time delay is governed by the dynamical friction timescale for the galaxies to lose orbital energy and momentum, and it depends on the mass ratios of the galaxies and the orbital parameters (\citealt{2008MNRAS.383...93B} and references therein). In addition to this difference in merger timescale, the growth in the stellar mass of a galaxy is not always proportional to the growth in its dark matter halo mass. A recent analysis of the galaxy catalogue in the Millennium simulation \citep{2007arXiv0708.1814G} finds galaxy growth via major mergers to depend strongly on stellar mass, where mergers are more important in the buildup of stellar masses in massive galaxies while star formation is more important in galaxies smaller than the Milky Way. Extending the analysis of this paper to the mergers of {\it subhaloes} in the Millennium simulation will provide the essential link between their and our results. For similar reasons, our results for the evolution of the dark matter halo merger rate per Gyr ($(1+z)^{n_m}$ with $n_m\sim2-2.3$) cannot be trivially connected to the observed merger rate of \emph{galaxies}. It is nonetheless interesting to note that a broad disagreement persists in the observational literature of galaxy merger rates. The reported power law indices $n_m$ have ranged from 0 to 5 (see, for example, \citealt{1994ApJ...429L..13B, 1994ApJ...435..540C, 1995ApJ...445...37Y, 1995ApJ...454...32W, 1997ApJ...475...29P, 2000MNRAS.311..565L, 2002ApJ...565..208P,2003AJ....126.1183C, 2004ApJ...601L.123B,2004ApJ...612..679L, 2004ApJ...617L...9L}). \cite{2006ApJ...652...56B} followed the redshift evolution of subhalo mergers in N-body simulations and provided a more detailed comparison with recent observations by, e.g., \cite{2004ApJ...617L...9L} that find $n_m<1$. They attributed such a weak redshift evolution in the number of close companions per galaxy to the fact that the high merger rate per halo at early times is counteracted by a decrease in the number of haloes massive enough to host a galaxy pair. The merger rates in this paper are global averages over all halo environments. The rich statistics in the Millennium simulation allow for an in-depth analysis of the environmental dependence of dark matter halo merger rates, which we will report in a subsequent paper (Fakhouri \& Ma, in preparation). | 7 | 10 | 0710.4567 |
0710 | 0710.1945_arXiv.txt | {The ARGO-YBJ experiment is a full coverage EAS-array installed at the YangBaJing Cosmic Ray Laboratory (4300 m a.s.l., Tibet, P.R. China). We present the results on the angular resolution measured with different methods with the full central carpet. The comparison of experimental results with MC simulations is discussed.} \begin{document} | The ARGO-YBJ detector is constituted by a single layer of Resistive Plate Chambers (RPCs). This carpet has a modular structure, the basic unit is a cluster, composed by 12 RPCs (2.8$\times$1.25 m$^2$ each). Each chamber is read by 80 strips, logically organized in 10 independent pads\cite{nim_argo}. The central carpet, constituted by 10$\times$13 clusters with $\sim$93$\%$ of active area, is enclosed by a guard-ring partially instrumented ($\sim$40$\%$) in order to improve rejection capability for external events. A lead converter 0.5 cm thick will uniformly cover the apparatus in order to improve the angular resolution. Since December 2004 the pointing accuracy of the detector has been studied, during the detector setting-up, with 3 different carpet areas: 42 clusters (ARGO-42, $\sim$1900 m$^2$), 104 clusters (ARGO-104, $\sim$4600 m$^2$) and the full central carpet, 130 clusters (ARGO-130, $\sim$5800 m$^2$), yet without any converter sheet. The data have been collected with a so-called {\it "Low Multiplicity Trigger"}, requiring at least $20$ fired pads on the whole detector. | Since December 2004 increasing fractions of ARGO-YBJ detector have been put in data taking even with a reduced duty-cycle due to installation and debugging operations. In this paper we presented a measurement of the pointing accuracy of the ARGO-130 detector. The capability of reconstructing the primary shower direction has been investigated with the chessboard method and with a preliminary Moon shadow analysis. Studies are in progress in order to determine the final angular resolution. | 7 | 10 | 0710.1945 |
0710 | 0710.5427_arXiv.txt | \footnotesize\ The efficient numerical solution of Non-LTE multilevel transfer problems requires the combination of highly convergent iterative schemes with fast and accurate formal solution methods of the radiative transfer (RT) equation. This contribution\footnote{Published in 1999 in the book {\it Solar Polarization}, edited by K.N. Nagendra \& J.O. Stenflo. Kluwer Academic Publishers, 1999. (Astrophysics and Space Science Library ; Vol. 243), p. 219-230} begins presenting a method for the formal solution of the RT equation in three-dimensional (3D) media with horizontal periodic boundary conditions. This formal solver is suitable for both, unpolarized and polarized 3D radiative transfer and it can be easily combined with the iterative schemes for solving non-LTE multilevel transfer problems that we have developed over the last few years. We demonstrate this by showing some schematic 3D multilevel calculations that illustrate the physical effects of horizontal radiative transfer. These Non-LTE calculations have been carried out with our code MUGA 3D, a 3D multilevel Non-LTE code based on the Gauss-Seidel iterative scheme that Trujillo Bueno and Fabiani Bendicho (1995) developed for RT applications. | To what extent can we trust diagnostic results obtained with the assumption that the solar atmospheric plasma is composed of {\it homogeneous} plane-parallel layers or via approximations that neglect {\it horizontal} radiative transfer (RT) effects? How important are the errors in the magnetic fields, temperatures and velocities inferred by confronting spectro-polarimetric observations with Non-LTE 1D RT model calculations? Clearly, to provide proper answers to questions like these requires to develop first efficient 3D RT methods that allow Non-LTE effects in complex atomic models with many levels and transitions to be rigorously investigated. There is a second reason which makes the development of fast iterative methods for 3D Non-LTE RT so relevant. This is because processes of energy exchange by radiation play an important role in the structure and dynamical behaviour of the stellar magnetized plasma. Thus, for instance, if one wishes to perform time-dependent radiation hydrodynamics simulations similar to those carried out by Carlsson and Stein (1997), but in 3D instead of 1D, it turns out to be imperative to have first access to numerical methods capable of accurately yielding the self-consistent atomic level populations at the cost of only {\it very} few formal solution times. The efficient solution of multilevel transfer problems requires the combination of a highly convergent iterative scheme with a fast formal solver of the RT equation. In Section 2 we briefly comment on a hierarchy of iterative schemes that can be applied for solving multilevel Non-LTE problems with increasing improvements in the convergence rate and total computational work. The 3D multilevel transfer calculations that we present in this contribution have been obtained by combining a highly convergent iterative scheme based on Gauss-Seidel iteration (Trujillo Bueno and Fabiani Bendicho, 1995) with a fast 3D formal solver that has parabolic accuracy (see Section 3). As was the case with our 2D formal solver, our generalization to 3D is based on the ``short-characteristics'' method of Kunasz and Auer (1988). Our 3D multilevel code is called MUGA 3D (``Multi-level Gauss-Seidel Method'') and it is substantially {\it faster} than our code MALI 3D, which is based on Jacobi iteration (see Rybicki and Hummer, 1991; Auer, Fabiani Bendicho and Trujillo Bueno, 1994). Section 3 briefly describes our 3D formal solver as applied to the scalar transfer equation for the specific intensity (I). In order to be able to consider 3D atmospheric models where solar plasma structures repeat themselves along the horizontal directions we choose horizontal periodic boundary conditions along the Cartesian coordinates X and Y. Although we do not give any details here, we have also generalized to 3D the Stokes-vector 1D formal solver method developed by Trujillo Bueno (1998), which is based on the matrix exponential approximation to the evolution operator. In Section 4 we show some illustrative 3D multilevel transfer calculations for a 5-level Ca II model atom where the H, K and infrared triplet lines are treated simultaneously, taking fully into account the {\it interlockings} by which photons are converted back and forth between the different line transitions in the assumed 3D medium. Here we consider schematic 3D solar models characterized by horizontal sinusoidal temperature inhomogeneities. With the help of these 3D multilevel calculations we are able to illustrate some subtle effects of horizontal radiative transfer that are important for the correct interpretation of high spatial resolution observations. Finally, Section 5 gives our conclusions. | We have developed a 3D multilevel code (MUGA-3D) that combines the Gauss-Seidel iterative scheme of Trujillo Bueno and Fabiani Bendicho (1995) with a 3D formal solver that uses horizontal periodic boundary conditions. With this new code we have performed some 3D multilevel simulations that highlight the importance of carefully investigating the effects of horizontal radiative transfer using realistic atmospheric and atomic models. We point out that our 3D formal solver can be used not only for solving ``unpolarized'' multilevel transfer problems, but also resonance line polarization and Hanle effect problems, like those considered in these Proceedings by Manso Sainz and Trujillo Bueno (1999), Paletou {\it et. al.} (1999) or Dittmann (1999). This is because, for these polarization transfer cases, the absorption matrix is diagonal. As a result, we have similar equations for the Stokes I, Q and U parameters. However, for the solution of more general polarization transfer problems, like the Non-LTE Zeeman line transfer case considered by Trujillo Bueno and Landi Degl'Innocenti (1996), but in 3D instead of 1D, one needs a 3D formal solution method of the Stokes-vector transfer equation. This is because here the absorption matrix turns out to be a full $4\times4$ matrix and all the Stokes parameters are coupled together. To this end we have generalized to 3D the Stokes-vector formal solver developed by Trujillo Bueno (1998), which can be considered as a generalization to polarization transfer of the short-characteristics method. We would like to end this paper by saying that over the last 10 years we have witnessed impressive progress concerning the development of highly convergent iterative schemes and accurate formal solvers for RT applications. Now it is time to apply them with physical intuition in order to improve our knowledge of the Sun, its magnetic field and its polarized spectrum. | 7 | 10 | 0710.5427 |
0710 | 0710.1338_arXiv.txt | The two-body problem in general relativity is reviewed, focusing on the final stages of the coalescence of the black holes as uncovered by recent successes in numerical solution of the field equations. | A black hole is one of the most fascinating and enigmatic predictions of Einstein's theory of general relativity. Its interior can have rich structure and is intrinsically dynamical, where space and time itself are inexorably led to a singular state. The exterior of an isolated black hole is, on the other hand, remarkably simple, described uniquely by the stationary Kerr solution. The dynamics of black holes are governed by laws analogous to the laws of thermodynamics, and indeed when quantum processes are included, emit Hawking radiation with a characteristic thermal spectrum. Most remarkable however, is that black holes, ``discovered'' purely through thought and the mathematical exploration of a theory far removed from every day experience, appear to be ubiquitous objects in our universe. The evidence that black holes exist, though circumstantial, is quite strong~\cite{Narayan:2005ie}. The high luminosity of quasars and other active galactic nuclei (AGN) can be explain by gravitational binding energy released through gas accretion onto supermassive ($10^6-10^9 \msun$) black holes at the centers of the galaxies~\cite{Rees:1984si,Ferrarese:2004qr}, several dozen X-ray binary systems discovered to date have compact members too massive to be neutron stars and exhibit phenomena consistent with matter interactions originating in the strong gravity regime of an inner accretion disk~\cite{McClintock:2003gx}, and the dynamical motion of stars and gas about the centers of nearby galaxies and our Milky Way Galaxy infer the presence of very massive, compact objects there, the most plausible explanation being supermassive black holes~\cite{Gebhardt:2000fk,Schodel:2002py,Ghez:2003qj}. To conclusively prove that black holes exist one needs to ``see'' them, or conversely see the compact objects masquerading as black holes. The only direct way of observing black holes is via the gravitational waves they emit when interacting with other matter/energy (an isolated black hole does not radiate). The quadrupole formula says that the typical magnitude $h$ of the gravitational waves emitted by a binary with reduced mass $\mu$ on a circular orbit measured a distance $r$ from the source is (for a review of gravitational wave theory see~\cite{Flanagan:2005yc}) \be h=\frac{16 \mu v^2}{r}, \ee where $v$ is the average tangential speed of the two members in the binary (and geometric units are used---Newton's constant $G=1$ and the speed of light $c=1$). This formula suggests that the strongest sources of gravitational waves are simply the most massive objects that move the fastest. To reach large velocities in orbit, the binary separation has to be small; black holes, being the most compact objects allowed in the theory, can reach the closest possible separations and hence largest orbital velocities. Therefore, modulo questions about source populations in the universe, a binary black hole interaction offers one of most promising venues of observing black holes through gravitational wave emission. Joseph Weber pioneered the science of gravitational wave detection with the construction of resonant bar detectors. Weber claimed to have detected gravitational waves~\cite{Weber:1969bz}, though no similar detectors constructed following his claims were able to observe the putative (or any other) source, and the general consensus is that given the sensitivity of Weber's detector and expected strengths of sources it is very unlikely that it was a true detection~\cite{Thorne:1980rt}. Note that the {\em existence} of gravitational waves is not in doubt---the observed spin down rate of the Hulse-Taylor binary pulsar~\cite{Hulse:1974eb} and several others discovered since, is in complete accord with the general relativistic prediction of spin down via gravitational wave emission. Today a new generation of gravitational wave detectors are operational, including laser interferometers (LIGO~\cite{LIGO}, VIRGO~\cite{VIRGO}, GEO600~\cite{GEO}, TAMA~\cite{TAMA}) and resonant bar detectors (NAUTILUS~\cite{NAUTILUS}, EXPLORER~\cite{EXPLORER}, AURIGA~\cite{AURIGA}, ALLEGRO~\cite{ALLEGRO}, NIOBE~\cite{NIOBE}). A future space-based observatory is planned (LISA~\cite{LISA}), and pulsar timing and cosmic microwave background polarization measurements also offer the promise of acting as gravitational wave ``detectors'' (for reviews see~\cite{Maggiore:1999vm,Cutler:2002me}). The ultimate success of gravitational wave detectors, in particular with regards to using them as more that simply detectors, but tools to observe and understand the universe, relies on source modeling to predict the structure of the waves emitted during some event. Even if an event is detected with a high signal-to-noise ratio (SNR), there simply is not enough information contained in such a one dimensional time series to ``invert'' it to reconstruct the event; rather template banks of theoretical waveforms from plausible sources need to be built and used to decode the signal. In rare cases an electromagnetic counterpart may be detected, for example during a binary neutron star merger if this is a source of short gamma ray bursts, which could identify the event without the need for a template. Though even in such a case, to extract information about the event, its environment, etc. requires source modeling. Gravitational wave detectors have therefore provided much of the impetus for trying to understand the nature of binary black hole collisions, and the gravitational waves emitted during the process. However, from a theoretical perspective black hole collisions are fantastic probes of the dynamical, strong-field regime of general relativity. What is already know about this regime---the inevitability of spacetime singularities in gravitational collapse via the singularity theorems of Penrose and Hawking~\cite{Hawking:1969sw,hawking_ellis}; the spacelike, chaotic ``mixmaster'' nature of these singularities conjectured by Belinsky, Khalatnikov and Lifshitz (BKL)~\cite{Berger:1998us}; the null, mass-inflation singularity discovered by Poisson and Israel~\cite{Poisson:1990eh} that, together with regions of BKL singularities could generically describe the interiors of black hole; the rather surprising discovery of critical phenomena in gravitational collapse by Choptuik~\cite{Choptuik:1992jv,Gundlach:1999cu}; etc---together with the sparsity of solutions (exact, numerical or perturbative), suggests there is potentially a vast landscape of undiscovered phenomena. Of particular interest, and potential application to high energy particle collision experiments, are ultra-relativistic black hole collisions. It is beyond the scope of this article to delve much into these aspects of black hole coalescence, though a brief overview of this will be given in Sec.~\ref{sec_he}. The two body problem in general relativity, introduced in more detail in Sec.~\ref{sec_2bdy}, is a very rich and complicated problem, with no known closed-form solution. Perturbative analytic techniques have been developed to deal with certain stages of the problem, in particular the inspiral prior to merger and ringdown after merger. Numerical solution of the full field equations are required during the merger, and this aspect of the problem is the main focus of this article. Much effort has been expended by the community over the past 15-20 years to numerically solve for merger spacetimes, and within the last two years an understanding of this phase of the two body problem is finally being attained. Sec.~\ref{sec_num} summaries the difficulties in discretizing the field equations, and describes the methods known at present that work for black hole collisions, namely {\em generalized harmonic coordinates} and {\em BSSN} (Baumgarte-Shapiro-Shibata-Nakamura) with moving punctures. Preliminary results are discussed in Sec.~\ref{sec_res}, though given the rapid pace at which the field is developing much of this will probably be dated in short order. Sec.~\ref{sec_imp} concludes with a discussion of some astrophysical and other implications of the results. | The two body problem in general relativity is a fascinating, rich problem that is just beginning to be fully revealed by recent breakthroughs in numerical relativity. At the same time, a new generation of gravitational wave detectors promise to offer us a view of the universe via the gravitational wave spectrum for the first time. Black hole mergers are a promising source for gravitational waves, and detecting them would provide direct evidence for these remarkable objects, while providing much information about their environments. Suggestions that there might be more than four spacetime dimensions offers the astonishing possibility that black holes could be produced by proton collisions at $TeV$ energies, which will be reached by the Large Hadron Collider, planned to begin operation within a year. Given all this, it is difficult not to be excited about what might be learnt about the universe from the smallest to largest scales during the next decade, and that black hole collisions could have something important to say at both extremes. \bigskip \noindent{\bf{\em Acknowledgments:}} I would like to thank Alessandra Buonanno, Matthew Choptuik, Gregory Cook, Charles Gammie, David Garfinkle, Steven Gubser, Carsten Gundlach, Luis Lehner, Jeremiah Ostriker, Don Page, David Spergel and Ulrich Sperhake for many stimulating conversations related to some of the discussion presented here. | 7 | 10 | 0710.1338 |
0710 | 0710.5388_arXiv.txt | Recently, it has been shown that the standard Nambu-Jona-Lasinio (NJL) model is not able to reproduce the correct QCD behavior of the gap equation at large density, and therefore a different cutoff procedure at large momenta has ben proposed. We found that, even with this density dependent cutoff procedure, the pure quark phase in neutron stars (NS) interiors is unstable, and we argue that this could be related to the lack of confinement in the original NJL model. | 7 | 10 | 0710.5388 |
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0710 | 0710.5207_arXiv.txt | If massive black holes (BHs) are ubiquitous in galaxies and galaxies experience multiple mergers during their cosmic assembly, then BH binaries should be common albeit temporary features of most galactic bulges. Observationally, the paucity of active BH pairs points toward binary lifetimes far shorter than the Hubble time, indicating rapid inspiral of the BHs down to the domain where gravitational waves lead to their coalescence. Here, we review a series of studies on the dynamics of massive BHs in gas-rich galaxy mergers that underscore the vital role played by a cool, gaseous component in promoting the {\it rapid formation of the BH binary}. The BH binary is found to reside at the center of a massive self-gravitating nuclear disc resulting from the collision of the two gaseous discs present in the mother galaxies. Hardening by gravitational torques against gas in this grand disc is found to continue down to sub-parsec scales. The eccentricity decreases with time to zero and when the binary is circular, accretion sets in around the two BHs. When this occurs, each BH is endowed with it own small-size ($\simlt 0.01$ pc) accretion disc comprising a few percent of the BH mass. Double AGN activity is expected to occur on an estimated timescale of $\simlt 1$ Myr. The double nuclear point--like sources that may appear have typical separation of $\simlt 10$ pc, and are likely to be embedded in the still ongoing starburst. We note that a potential threat of binary stalling, in a gaseous environment, may come from radiation and/or mechanical energy injections by the BHs. Only short--lived or sub--Eddington accretion episodes can guarantee the persistence of a dense cool gas structure around the binary necessary for continuing BH inspiral. | Dormant black holes (BHs) with masses in excess of $\simgt 10^6\msun$ are ubiquitous in bright galaxies today (Kormendy \& Richstone 1995; Richstone 1998). Relic of an earlier active phase as quasars, these massive BHs appear a clear manifestation of the cosmic assembly of galaxies. The striking correlations observed between the BH masses and properties of the underlying hosts (Magorrian \etal 1998; Ferrarese \& Merritt 2000; Gebhardt et al. 2000; Graham \& Driver 2007) indicates unambiguously that BHs evolve in symbiosis with galaxies, affecting the environment on large--scales and self-regulating their growth (Silk \& Rees 1998; King 2000; Granato et al. 2004; Di Matteo, Springel \& Hernquist, 2005). According to the current paradigm of structure formation, galaxies often interact and collide as their dark matter halos assemble in a hierarchical fashion (Springel, Frenk \& White 2006), and BHs incorporated through mergers into larger and larger systems are expected to evolve concordantly (Volonteri, Haardt \& Madau 2003). In this astrophysical context, close BH {\it pairs} form as natural outcome of binary galaxy mergers (Kazantzidis et al. 2005). In our local universe, one outstanding example is the case of the ultra--luminous infrared galaxy NGC 6240, an ongoing merger between two gas--rich galaxies (Komossa \etal 2003; for a review on binary black holes see also Komossa 2006). {\it Chandra} images have revealed the occurrence of two nuclear X-ray sources, 1.4 kpc apart, whose spectral distribution is consistent with being two accreting massive BHs embedded in the diluted emission of a starburst. Similarly, Arp 299 (Della Ceca et al. 2002; Ballo et al. 2004) is an interacting system hosting an obscured active nucleus, and possibly a second less luminous one, distant several kpc away. A third example is the elliptical galaxy 0402+369 where the cleanest case of a massive BH {\it binary} has been recently discovered. Two compact variable, flat--spectrum active nuclei are seen at a projected separation of only 7.3 pc (Rodriguez et al. 2006). Arp 299, NGC 6240, and 0402+369 may just highlight different stages of the BH dynamical evolution along the course of a merger, with 0402+369 being the latest, most evolved phase (possibly related to a dry merger). Energy and angular momentum losses due to gravitational waves are not yet significant in 0402+392, so that stellar interactions and/or material and gravitational torques are still necessary to bring the BHs down to the domain controlled by General Relativity. From the above considerations and observational findings, it is clear that binary BH inspiral down to coalescence is a major astrophysical process that can occur in galaxies. It is accompanied by a gravitational wave burst so powerful to be detectable out to very large redshifts with current planned experiments like the Laser Interferometer Space Antenna ({\it LISA}; Bender et al. 1994; Vitale et al. 2002). These extraordinary events will provide not only a firm test of General Relativity but also a view, albeit indirect, of galaxy clustering (Haehnelt 1994; Jaffe \& Backer 2003; Sesana et al. 2005). With {\it LISA}, BH masses and spins will be measured with such an accuracy (Vecchio 2004) that it will be possible to trace the BH mass growth across all epochs. Interestingly, {\it LISA} will explore a mass range between $10^3\msun$ and $10^7\msun$ that is complementary to that probed by the distant massive quasars ($>10^7\msun$), providing a complete census of the BHs in the universe. Both minor as well as major mergers with BHs accompany galaxy evolution in environments that involve either gas-rich (wet) as well as gas-poor (dry) galaxies. Thus, the dynamical response of galaxies to BH pairing should differ in many ways according to their properties. Exploring the expected diversities in a self-consistent cosmological scenario is a major challenge and only recently, with the help of high-resolution N-body/SPH simulations, it has become possible to ``start'' addressing a number of compelling issues. Galaxy mergers cover cosmological volumes (a few to hundred kpc aside), whereas BH mergers probe volumes of only few astronomical units or less. Thus, tracing the BH dynamics with scrutiny requires N-Body/SPH force resolution simulations spanning more than nine orders of magnitude in length. For this reason, two complementary approaches have been followed in the literature. A statistical approach (based either on Monte Carlo realizations of merger trees or on N-Body/SPH large scale simulations) follows the collective growth of BHs inside dark matter halos. Supplemented by semi-analytical modeling of BH dynamics (Volonteri et al. 2003) or/and by sub-grid resolution criteria for accretion and feedback (Springel \& Hernquist 2003; Springel, Di Matteo \& Hernquist 2005), these studies have proved to be powerful in providing estimates of the expected coalescence rates, and in tracing the overall cosmic evolution of BHs including their feedback on the galactic environment (Di Matteo, Springel \& Hernquist 2005; Di Matteo \etal 2007). The second approach, that we have been following, looks at individual binary collisions, as it aims at exploiting in detail the BH dynamics and some bulk physics from the galactic scale down to and within the BH sphere of influence. Both approaches, the collective and the individual, are necessary and complementary, the main challenge being the implementation of realistic input physics in the dynamically active environment of a merger. Following a merger, how can BHs reach the gravitational wave inspiral regime? The overall scenario was first outlined by Begelman, Blandford \& Rees (1980) in their seminal study on the dynamical evolution of BH pairs in pure stellar systems. They indicated three main roots for the loss of orbital energy and angular momentum: (I) dynamical friction against the stellar background acting on each individual BH; (II) hardening via 3--body scatterings off single stars when the BH binary forms; (III) gravitational wave back--reaction. Early studies explored phase (I) simulating the collisionless merger of spherical halos (Makino \& Ebisuzaki 1996; Milosavljevi\'c \& Merritt 2001; Makino \& Funato 2004). Governato, Colpi \& Maraschi (1994) in particular first noticed that when two equal mass halos merge, the twin BHs nested inside the nuclei are dragged effectively toward the center of the remnant galaxy by dynamical friction and form a close pair, but that the situation reverses in unequal mass mergers, where the less massive halo tidally disrupted leaves its ``naked'' BH wandering in the outskirts of the main halo. Thus, depending on the halo mass ratio and internal structure, the transition from phase (I) to phase (II) can be prematurely aborted or drastically relented. Similarly, the transit from phase (II) to phase (III) is not always secured, as the stellar content inside the ``loss cone" may not be rapidly refilled with fresh low--angular momentum stars to harden the binary down to separations where gravitational wave driven inspiral sets in (see, e.g., Milosavljevic \& Merritt 2001; Yu 2002; Berczik, Merritt \& Spurzem 2005; Sesana, Haardt \& Madau 2007). For an updated review on the last parsec problem and its possible solution (see Merritt 2006a; Gualandris \& Merritt 2007). Since BH coalescences are likely to be events associated with mergers of (pre--)galactic structures at high redshifts, it is likely that their dynamics occurred in gas dominated backgrounds, NGC 6240 being just the most outstanding case visible in our local universe. Other processes of BH binary hardening are expected to operate in presence of a dissipative gaseous component that we will highlight and study here. Kazantzidis \etal (2005) first explored the effect of gaseous dissipation in mergers between gas--rich disc galaxies with central BHs, using high resolution N--Body/SPH simulations. They found that the merger triggers large--scale gas dynamical instabilities that lead to the gathering of cool gas deep in the potential well of the interacting galaxies. In minor mergers, this fact is essential in order to bring the BHs to closer and closer distances before the less massive galaxy, tidally disrupted, is incorporated in the main galaxy. Moreover, the interplay between strong gas inflows and star formation leads naturally to the formation, around the two BHs, of a grand, massive ($\sim 10^9\msun$) gaseous disc on a scale smaller than $\sim 100$ pc. It is in this equilibrium circum--nuclear disc that the dynamical evolution of the BHs continues, after the merger has been subsided. Escala \etal (2005, hereinafter ELCM05; see also Escala \etal 2004) have been the first to study the role played by gas in affecting the dynamics of massive ($\sim 10^8\msun$) twin BHs in equilibrium Mestel discs of varying clumpiness. In both these approaches (i.e., in the large scale simulations of Kazantzidis et al., and in the equilibrium disc models of ELCM05) it was clear that the gas temperature is a key physical parameter and that a hot gas brakes the BHs inefficiently. Instead, when the gas is allowed to cool, the drag becomes efficient: the large enhancement of the local gas density relative to the stellar one leads to the formation of prominent density wakes that are decelerating the BHs down to the scale where they form a ``close'' binary. Later, binary hardening occurs under mechanisms that are only partially explored, and that are now subject of intense investigation. The presence of a cool circum--binary disc and of small--scale discs around each individual BH appear to be critical for their evolution down to the domain of gravitational waves driven inspiral. In this context there is no clear ``stalling problem'' that emerges from current hydrodynamical simulations but this critical phase need a more through, coherent analysis. The works by Kazantzidis \etal (2005) and ELCM05 have provided our main motivation to study the process of BH pairing along two lines: In gas-rich binary mergers, line (1) aims at studying the transit from state (A) of pairing when each BH moves individually inside the time-varying potential of the colliding galaxies, to state (B) when the two BHs dynamically couple their motion to form a binary. The transit from (A) $\to$ (B) requires exploring a dynamic range of five orders of magnitude in length from the cosmic scale of a galaxy merger of 100 kpc down to the parsec scale for BHs of million solar masses (i.e., BHs in the LISA sensitivity domain). After all transient inflows have subsided and a new galaxy has formed, the BH binary is expected to enter phase (C) where it hardens under the action of gas-dynamical and gravitational torques. Research line (2) aims at studying the braking of the BH binary from (B) $\to$ (C) and further in, exploring the possibility that during phase (C) two discs form and grow around each individual BH. As first discussed by Gould \& Rix (2000) the binary may later enter a new phase (D) controlled by the balance of viscous and gravitational torques in a circum--binary disc surrounding the BHs, in a manner analogous to the migration of planets in circum-stellar discs (a scenario particularly appealing when the BH mass ratio is less than unity). Phase (D) likely evolves into (E) when gravitational wave inspiral terminates the BH binary evolution. There is a number of key questions to address: \noindent (i) How does transition from state (A) $\to$ (B) depend on the gas thermodynamics? How do BHs bind? \noindent (ii) In the grand nuclear disc inside the remnant galaxy, how do eccentric orbits evolve? Do they become circular or highly eccentric? \noindent (iii) During the hardening through phase (B) and (C), do the BHs collect substantial amounts of gas to form cool individual discs? \noindent (iv) Can viscous torques drive the binary into the gravitational wave decaying phase? \noindent (v) Is there a threat of a {\it stalling} problem when transiting from (C) $\to$ (D) or from (D) $\to$ (E)? And, for which mass ratios and ambient conditions? | 7 | 10 | 0710.5207 |
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0710 | 0710.0368_arXiv.txt | We compute opacities for the electronic molecular band systems \ACrH -- \XCrH of CrH and CrD, and \AMgH -- \XMgH of MgH and MgD. The opacities are computed by making use of existing spectroscopic constants for MgH and CrH. These constants are adjusted for the different reduced masses of MgD and CrD. Frank-Condon factors are used to provide intensities for the individual vibronic bands. These results are used in the computation of synthetic spectra between \Tef = 1800 and 1200 K with an emphasis on the realisation of ``deuterium test'', first proposed by Bejar et al. (1999) to distinguish brown dwarfs from planetary mass objects. We discuss the possible use of CrD and MgD electronic bands for the ``deuterium test". We find CrD to be the more promising of the two deuterides, potentially, the most useful bands of CrH/CrD are the $\Delta v = +1$ and $\Delta v = -1$ at 0.795 and 0.968 \mum . | The ``deuterium test'' was suggested as a method of identifying planetary mass objects among cool objects (\Bejar et al. 1999, Chabrier et al. 2000). In practise, it was proposed to search for absorption lines of molecules containing deuterium (HDO, CrD, FeD, etc.). Deuterium is burnt in stellar interiors via the fusion reaction $^2$D(p,$\gamma$)$^3$He at temperatures ($T > 8 \times 10^5$~ K). The interiors of substellar objects with M $<$ 13~ \MJ, where \MJ is a Jovian mass (0.001 \Msun), do not reach temperatures high enough for deuterium to ignite (Saumon et al. 1996). As a result the deuterium abundance in the atmospheres of these objects is unchanged from the formation of these objects. This gives rise to the definition of a brown dwarf as an object which has insufficient mass to fuse $^1$H to $^4$He, but has sufficient mass to fuse $^2$H to $^3$He. By comparison, a planetary mass object has insufficient mass to ignite fusion of any sort (Saumon et al. 1996). In higher-mass objects such as stars, deuterium burning is completed comparatively quickly (t = 1 -- 3 million years) during the evolution prior to the star's period on the main sequence (D'Antona \& Mazitelli 1998). The deuterium depletion rate depends on the mass of the star or brown dwarf and thus the ``deuterium test'' can be used in a number of ways. \begin{itemize} \item Discern planets from brown dwarfs in a population of low-mass objects. \item Determine the evolutionary status of young objects in open clusters with ages of several million years. \item Study the evolution of the abundance of deuterium in the atmospheres of low-mass substellar objects, a phenomenon which is poorly understood. The rate of depletion of deuterium depends upon rotation, magnetic field strength, and other parameters which affect the efficiency of convection in low mass objects. \end{itemize} Such investigations can usefully be combined with the ``lithium test'' proposed by Rebolo et al. (1992), \Magazzu et al. (1993). For stars (M $>$ 75 \MJ), the burning of lithium, Li (p,$\alpha) ^4$He, becomes efficient at evolutionary stages preceding the main sequence at interior temperatures of T $>$ $2.5\times 10^6$ K (D'Antona \& Mazitelli 1998). The ``lithium test'' has been successfully applied to identify brown dwarfs in a population of ultracool dwarf stars (Rebolo et al. 1996, Basri 2000). The lithium test is relatively easily applied to M-dwarfs. The reason for this is that the resonance lines of neutral atomic lithium lie in the optical region of the spectrum, at 0.6708 \mum. One blemish with the ``lithium test'' is the severe blending of lithium lines with the background of TiO lines, nonetheless Li lines break through the molecular background. In the spectra of cooler M and L dwarfs Ti atoms are depleted onto dust particles (Tsuji et al. 1996, Jones \& Tsuji 1997, Pavlenko 1998) so TiO absorption is weakened. In the L dwarf regime the appearance of lithium lines contends with dust opacity and the wings of K I and Na I resonance lines (see Pavlenko et al. 2000 and references therein). The realisation of the ``deuterium test'' is considerably more difficult. Conventionally, the deuterium abundance of hot objects is determined from analysis of multicomponent features on the background of the $L_\alpha$ 0.1215 \mum ~H~I line seen in emission. Unfortunately, this method cannot be used in the case of ultracool dwarfs which are covered by a thick envelope of neutral hydrogen. Observations also yield H~I emission lines ($H_\alpha$ 0.6563 \mum) formed in the outermost layers of hot chromospheres or accretion disks around young stars. However, such a plasma is variable and may be polluted by interstellar material. Taking into account these circumstances, it is logical to analyse the spectral lines of deuterated molecules formed in the photospheric layers of ultracool dwarfs. The first investigation and analysis of the combined spectra of \HHO /HDO were carried out by Chabrier et al. (2000) and Pavlenko (2002). Due to the change in mass and the breakdown of molecular symmetry the vibration-rotation bands of HDO in the mid-infrared spectrum shift with respect to the \HHO bands. There are a few spectral regions which can be used for detection of HDO lines in the IR spectra of ultracool dwarfs: 3.5 -- 4, $\sim$5, 6-7 \mum (Pavlenko 2002). The main problem is that despite the shift in wavelength such HDO lines will be on a background of far stronger \HHO lines. A possible alternative to HDO are the diatomic hydrides. Strong molecular bands of diatomic metal hydrides such as MgH and CrH can be observed in the optical spectrum of ultracool dwarf stars. The MgH band system \AMgH - \XMgH can be observed at 0.47--0.6 \mum, and the CrH band system \ACrH -- \XCrH\ show absorption features at 0.6--1.5 \mum. The molecules MgH and CrH have been known in astrophysics for a long time. MgH has been more extensively studied than CrH, because it can be observed in the spectra of G to M type stars. The dissociation energy of MgH is very low (1.285 eV) so lines of this molecule are very sensitive to the temperature and gravity variations in stellar atmospheres. MgH lines were used to determine temperatures in the atmospheres of cool giants (Wyller 1961) and the Sun (Sinha et al. 1979, Sinha \& Joshi 1982), and for the determination of the surface gravity of stars (Bell \& Gustaffson 1981, Bell at al. 1985, Berdyugina \& Savanov 1992, Bonnel \& Bell 1993). The pure rotational spectrum of MgH and MgD radicals (\XMgH) in their ground state $v$=0 and $v$=1 vibrational modes has been studied by Ziurys et al. (1993). The first MgH linelist was computed by Kurucz (1993). Recently, more extensive studies of MgH transitions were performed by Weck et al. (2003, 2003a, 2003b) and Skory et al. (2003). Although CrH has been known since Gaydon \& Pearse (1937), its electronic spectrum remained relatively unstudied for many years. Engvold et al. (1980) identified lines of CrH in a spectrum of sunspots as formed by \XCrH -- \XCrH transitions. They used the results of studies of multiplicity of $\Sigma$-terms of CrH by Klehman \& Uhler (1959) and O'Connor (1967). Later Ram, Jarman \& Bernath (1993) performed a rotational analysis of 0--0 band of the \ACrH -- \XCrH electronic transition and obtained improved rotational constants for the $v'$ =0 vibrational state. Combining these results with those of Bauschlicher et al. (2001) and Lipus et al. (1991) for the vibrationally excited transitions, Burrows et al.(2002) computed an extended linelist for CrH. Recently Shin et al. (2005) have measured radiative lifetimes of the $v$ = 0,1 levels of \ACrH state of CrH. These measured lifetimes are about 16\% -- 45 \% longer that those obtained by Burrows at al. (2002). These results provide evidence that the oscillator strengths of Burrows et al. (2002) should be corrected by a factor of 0.8 for at least the transitions to the $v'$ = 0. The submillimeter spectra of CrH and CrD formed by pure rotational transitions in the ground electronic state, have been observed in the laboratory by Halfen \& Ziurys (2004). Electronic bands of MgD and CrD are likely to be located in the same spectral regions as the corresponding bands of MgH and CrH. In this paper we model the bands of these molecules to analyse the possibility of their use for the determination of the D/H ratio in the atmospheres of ultracool dwarfs. In section 2 we present a description of the procedures used to compute the molecular bands of CrH, CrD, MgH and MgD. In section 3 we present the vibrational-rotational constants of MgH, MgD, CrH and CrD. In section 4 we present the results of the computation of molecular bands. In section 5 we discuss the possibility of using the electronic bands of diatomic molecules for a realisation of the ``deuterium test''. | The detection of deuterated molecules in the spectra of ultracool dwarfs provide a challenge for both theoreticians and observers. Indeed, in the atmospheres of planetary mass objects (M $<$ 13 M$_J$) we cannot expect ratio D/H $> 2\times10^{-5}$. This means the lines of deuterated molecules should be about 5000 times weaker than those of the hydrides. The ideal case would be a spectral region where the molecular bands are not blended. So a crucial requirement is a large difference in the wavelengths of the band heads of the hydrides and deuterides. CrH appears to be more useful than MgH in the search for deuterated species. The band heads of CrD are displaced significantly from the band heads of CrH. Bands of CrH are observed in the spectra of the latest L dwarfs (Kirkpatrick et al. 1999). The CrD bands are located in the ``near infrared'' spectra, where fluxes are much higher than in the ``optical'' spectral regions. In this paper we show that the most useful bands for the realisation of the deuterium test are $\Delta v = +1$ and $\Delta v = -1$ ($\lambda$ = 0.795 and 0.968 \mum, respectively). The $\Delta v = -1$ band looks especially promising. It is located in the near infrared region with an absence of strong background absorption features. However, portions of this CrD band will be swamped by the 0-0 FeH Wing-Ford band at 0.99$\mu$m and possibly by water bands. Still, the case for CrD looks better than for HDO lines which are formed on a background of strong H2O lines (Pavlenko 2002). High quality line lists are required to test these possibilities fully. It is worth noting, that a potential problem lies in the possibility that Cr atoms are absorbed onto dust particles. The depletion of Cr will reduce the strength of both CrD and CrH bands. Fortunately as CrH bands are located in the same spectral region, we can ``scale'' the CrD depletion processes by adjusting the strength of CrH bands. For more precise studies, more accurate and detailed linelists of CrH and CrD are required. The calculation of such linelists would require new improved computations supported by new laboratory measurements. One problem that concerns us is that even once we have a good agreement with the model and experimental data for the MgH or CrH molecule, there are still perturbations about which we know very little from experiments done so far. Nonetheless, MgH is a non-starter for the deuterium test and as we note above, the model adopted in our paper is more likely to be reliable for CrH and CrD. It is worth noting that the use of the pure rotational-vibrational bands located in the mid and far infrared spectral region may offer an alternative. Indeed, the displacement between CrH and CrD rotation-vibration bands is even larger, than for the case of electronic bands. Nevertheless, we cannot be certain that we have identified the best candidate systems for the deuterium test. Future investigations of deuterated molecules in different spectral regions are important to determine which offers the best possibility for the realisation of this test. The ideal solution would be to detect lines of deuterated molecule(s) in different spectral regions. This presents a serious observational challenge which can only be met in combination with careful laboratory measurement and the computation of high quality molecular spectra. | 7 | 10 | 0710.0368 |
0710 | 0710.1472_arXiv.txt | Very recently, J~1128+5925 was found to show strong intraday variability at radio wavelengths and may be a new source with annual modulation of the timescale of its radio variability. Therefore, its radio variability can be best explained via interstellar scintillation. Here we present the properties of its optical variability for the first time after a monitoring program in 2007 May. Our observations indicate that in this period J~1128+5925 only showed trivial optical variability on internight timescale, and did not show any clear intranight variability. This behavior is quite different from its strong radio intraday variability. Either this object was in a quiescent state in optical in this period, or it is intrinsically not so active in optical as it is in radio regimes. | Blazars are the most variable subset of AGN. They show a variety of variability timescales. The longest timescales can be far longer than one year, while the shortest may be less than one hour. The variability with a timescale less than one day is often called intraday variability or IDV, as first reported by \citet{heeschen84}, \citet{witzel86}, and \citet{heeschen87}. Strong IDV phenomena have been observed in the radio domain in a large number of blazars. If interpreted as being source intrinsic, the short-timescale variability would require a very small emitting region and hence a very high apparent brightness temperature of $10^{16}\sim10^{21}\,\rm{K}$, which is far beyond the inverse-Compton limit of about $10^{12}\,\rm{K}$ \citep{keller69}. Alternatively, the IDV can be explained via extrinsic origin, e.g., via interstellar scintillation (ISS). A strong support to the ISS origin is the so-called annual modulation of the variability timescale, which is the result of the annual changes of the relative velocity vector between the scattering screen and the Earth as the Earth orbits around the Sun \citep[e.g.,][]{dennett02,dennett03}. Such annual modulation has been observed in a few IDV sources, as mentioned by \citet{gabanyi07}. Very recently, the flat-spectrum radio quasar J~1128+5925 was found to show strong IDV at centimeter wavelengths, and its IDV timescale displays an annual modulation \citep{gabanyi07}. Therefore, its IDV may be caused by ISS. In optical, there is no previous report on its variability. In order to know the properties of its optical variability and to make a comparison to those of its radio variability, we performed an optical monitoring program on this object in 2007 May. Here we present the results. | We performed an optical monitoring program on J~1128+5925 in the $R$-band from 2007 May 5 to 29. Our monitoring results indicate that in this period J~1128+5925 only showed trivial optical variability on internight timescale, and did not show any clear intranight variability. Either this object was in a quiescent state in optical regimes in this period, or it is intrinsically not as active at optical as it is at radio wavelengths. Some blazars that show strong IDV in radio regimes also display rapid and strong variability in optical regimes, such as S5~0716+714 \citep[e.g.,][] {wu05,wu07,montagni06,pollock07}. This doesn't seem to be the case for J~1128+5925. This object exhibits strong IDV at radio wavelengths, but not at optical wavelengths. It is easy to explain this difference if the optical and radio variabilities come from different origins: The optical variability may be intrinsic to the source (the ISS cannot change the optical flux), while the radio variability is mainly the result of ISS, as implied by the observations of \citet{gabanyi07}. We present the first report on the optical variability of J~1128+5925 in this paper. However, because of the solar conjunction and observations of other targets with the telescope, our monitoring did not last long. More observations are needed to know whether or not this object is always optically quiescent. Multi-band optical monitoring is also necessary to constrain its optical variability in more detail. Of particular interest is to carry out simultaneous optical and radio monitorings on this object in order to make a more direct comparison between the variabilities at these two wavelengths. Future campaigns can investigate whether there is correlated optical IDV when strong radio IDV is observed. If such correlations are detected, it would be strong evidence that both the optical and radio variability structures are intrinsic to the source, as in the case of S5~0716+714 \citep{quirren91,wagner95,wagner96}. The broadband variability is also helpful to derive for this object some basic parameters, such as the mass of the central supermassive black hole, the boosting factor of the relativistic jet, etc \citep[e.g.,][]{fan05}. | 7 | 10 | 0710.1472 |
0710 | 0710.4582_arXiv.txt | We compare the stellar structure of the isolated, Local Group dwarf galaxy Pegasus (DDO\,216) with low resolution HI maps from \cite{young2003}. Our comparison reveals that Pegasus displays the characteristic morphology of ram pressure stripping; in particular, the HI has a ``cometary'' appearance which is not reflected in the regular, elliptical distribution of the stars. This is the first time this phenomenon has been observed in an isolated Local Group galaxy. The density of the medium required to ram pressure strip Pegasus is at least $10^{-5} - 10^{-6}$\,cm$^{-3}$. We conclude that this is strong evidence for an inter-galactic medium associated with the Local Group. | \cite{einasto1974} first highlighted that dwarf satellites of large galaxies tend to be gas deficient compared to isolated dwarfs. The former generally have little or no ongoing star formation and the stars are pressure supported (dwarf spheroidal, dSph). The latter generally have ongoing star formation and the gas dynamics show that rotational support is important (dwarf irregular, dIrr). ``Transition'' dwarfs are gas-rich and, unlike dIrr galaxies, have little or no detectable HII regions, although they usually show indications of recent star formation. The processes by which dwarf galaxies loose their gas are not fully understood. Internal feedback, particularly winds from supernovae, are likely important (\citealt{dekel1986}) and the existence of the position-morphology relation clearly indicates that environmental influences are significant. \cite{mayer2006} show that it is possible for dwarf galaxies to be ram pressure stripped of some of their gaseous component in a hot halo of the Milky Way or M31. This idea was originally proposed by \cite{lin1983}, who calculated the density of the medium required to be of order $10^{-6}$\,cm$^{-3}$. There have been no direct detections of such a medium, although recently \cite{nicastro2002,nicastro2003} and \cite{sembach2003} have detected OVI absorption which they attribute to hot gas associated with either a Milky Way corona or a Local Group medium. In this {\it Letter}, we compare the stellar and gaseous structure of the isolated, transition-type, dwarf galaxy Pegasus (DDO216). We show that it displays the characteristic signature of ram pressure stripping and conclude that this is strong evidence for hot gas associated with the Local Group. Table~1 summarises some of the observed properties of Pegasus. We adopt the distance estimate by \cite{mcconnachie2005a}, $D \simeq 919$\,kpc, derived from the same photometry used in this Letter. \begin{table}[htdp] \begin{center} \caption{Summary of observed parameters for the Pegasus (DDO216) dwarf galaxy} \begin{tabular}{lll} Parameter & Value & Reference \\ \hline $\alpha$ (J2000) & 23h 28m 36.2s& --- \\ $\delta$ (J2000) & +14$^\circ$ 44$^\prime$ 35$^{\prime\prime}$& --- \\ $(l, b)$ & $(94.8^\circ, -43.6^\circ)$ & --- \\ $M_V~{\it(L_V)}$ & -12.9~$(1.24 \times 10^7\,L_\odot)$ & \cite{mateo1998a} \\ $M_{HI}$ & $4.06 \times 10^6\,M_\odot$ & \cite{young2003}\\ $v_\odot$ & -183\,km\,s$^{-1}$ & \cite{young2003} \\ $v_r/\sigma$ & $1.7$ & \cite{mateo1998a} \\ Distance & 24.82 $\pm$ 0.07 (919\,kpc) & \cite{mcconnachie2005a} \\ & 24.4 $\pm$ 0.2 & \cite{gallagher1998} \\ & 24.9 $\pm$ 0.1 & \cite{aparicio1994} \\ \hline \end{tabular} \end{center} \label{distances} \end{table} | \subsection{Comparison of stellar and HI contours} The top-right panel of Figure~1 shows the tangent-plane projection of the spatial distribution of objects identified as stellar from our INT~WFC observations of Pegasus. The dotted lines in this panel (and the remaining panels of Figure~1) correspond to the approximate edges of each CCD of the INT~WFC. Only objects which lie within $1-\sigma$ of the stellar locus in both the $V^\prime$ and $i^\prime-$band observations are shown. The hole at the center of the main body of Pegasus is due to severe crowding which causes incompleteness. The bottom-left panel of Figure~1 shows a contour map of the density distribution of stars. The first contour is $2-\sigma$ above the background, and the contours correspond to $2.2, 5.0, 8.6, 13.2, 19,0, 26.3, 35.7, 47.5$ and $62.5$\,stars\,arcmin$^{-2}$. The contour map was made in the standard way and follows exactly the methodology described in \cite{mcconnachie2006b}. This panel shows that Pegasus is significantly more extended than suggested by the image in the first panel. The bottom-right panel of Figure~1 shows the stellar density distribution as a grey-scale with square-root scaling. The red contours are the low-resolution HI distribution from \cite{young2003}. The contours correspond to column densities of $0.1, 0.2, 0.4, 0.8, 1.6, 3.2, 6.4$ and $12.8 \times 10^{20}$\,cm$^{-2}$. Whereas the stars are distributed in a regular ellipse (typical of a flattened spheroid or an inclined disk) the HI has a ``cometary'' appearance; the contours in the south-east are more closely packed and do not extend as far as in the north-west. \subsection{Is Pegasus being ram pressure stripped?} \begin{figure*} \begin{center} \includegraphics[angle=270, width=8.9cm]{f1a.eps} \includegraphics[angle=270, width=8.9cm]{f1b.ps} \includegraphics[angle=270, width=8.9cm]{f1c.ps} \includegraphics[angle=270, width=8.9cm]{f1d.ps} \caption{Projections in the tangent plane $\left(\xi, \eta\right)$ of the structure of the Pegasus dwarf galaxy (DDO216) with the orientation of the field indicated. The dotted lines trace the approximate edges of the four CCDs of the INT~WFC. Top-left panel: The reduced $V^\prime-$band image of Pegasus taken with the INT~WFC. Also marked are the positions of fields analysed in previous studies of this galaxy; the fields analysed by \cite{gallagher1998} are shown in green as the largest rectangular field and the small WFPC2 footprint, the field analysed by \cite{aparicio1994} is shown in blue as the smallest rectangular field, and the field studied by \cite{hoessel1982} is shown in red as the medium sized rectangle. Top-right panel: The distribution of all sources confidently identified as stellar in both the $V^\prime$ and $i^\prime-$bands. The hole at the center of Pegasus is due to severe crowding causing the photometry to become seriously incomplete. Bottom-left panel: The stellar density distribution of Pegasus shown as a contour map. The first contour is $2-\sigma$ above the background and the 9 contours correspond to $2.2, 5.0, 8.6, 13.2, 19,0, 26.3, 35.7, 47.5$ and $62.5$\,stars\,arcmin$^{-2}$. Bottom-right panel: the stellar density distribution is shown as a grey-scale with square-root scaling. The red contours show the low resultion HI distribution from \cite{young2003}. Contours correspond to $0.1, 0.2, 0.4, 0.8, 1.6, 3.2, 6.4$ and $12.8 \times 10^{20}$\,cm$^{-2}$. Also shown are the projected directions to all galaxies within $\sim 500$\,kpc of Pegasus which have a significant gaseous content. Pegasus displays the characteristic morphology of ram pressure stripping.} \label{fig} \end{center} \end{figure*} The shape of the low resolution HI contours in Pegasus - the smooth, compressed contours in the south-east and the ``tail'' to the north-west - is very similar to the simulated morphology of gas undergoing ram pressure striping (e.g., \citealt{stevens1999,mori2000,marcolini2003,roediger2005,mayer2006}). Observationally, the M81 group dwarf galaxy Holmberg~II is observed to have a similar morphology (\citealt{bureau2002}), interpreted as evidence of an intra-group medium. In clusters of galaxies, ram pressure stripping of galaxies by an intra-cluster medium is used to explain various observations, including the deficit of HI in cluster spiral galaxies compared to field spirals (e.g., \citealt{giovanelli1985}). Indeed, several individual galaxies in the Virgo Cluster have been shown to display gaseous morphologies indicative of ram-pressure stripping (\citealt{vollmer2000,vollmer2004,vollmer2005}). What else could explain the peculiar appearance of Pegasus? While tidal stripping by large galaxies can affect the structure of dwarf galaxies (e.g., \citealt{penarrubia2007b}), the closest large galaxy to Pegasus is M31 at $\sim 470$\,kpc (all the distance estimates in Table~1 place Pegasus at $> 400\,$kpc from M31). Even if we assume Pegasus is a weakly-bound satellite of M31, tidal effects at this distance are minimal. If Pegasus was disrupted at pericenter, it is unlikely that the gas would still show signs of current disturbance. Further, tidal stripping tends to produce symmetrical distortions and both gas and stars should be affected. However, these are inconsistent with the structure of Pegasus that we observe. Could the appearance of Pegasus be due to internal effects rather than external influences? Enhanced star formation in the south-east of Pegasus could produce winds which remove gas from this region. However, if this is the case then the densely packed contours in the south-east should have a more concave, rather than convex, shape. For example, \cite{young2007} discuss a gas cloud associated with the Phoenix dwarf galaxy and conclude that it was blown out by supernovae winds based in part on the concave shape of its contours. An alternative explanation for the HI morphology of Pegasus is that it consists of multiple HI clouds, the sum total of which has a cometary appearance. Figure~6 of \cite{young2003} is a position-velocity diagram of Pegasus along its major axis. It shows a gradient in velocity and two main concentrations of HI which \cite{young2003} interpret as two distinct HI clouds. The strength of the secondary feature ($v \sim -200$\,km\,s$^{-1}$) is weaker than the main feature ($v \sim -180$\,km\,s$^{-1}$) and they join at relatively high column density (between the $8 - 16 - \sigma$ contour levels). An alternative explanation of the data is that the overall velocity gradient is a result of ram-pressure stripping. The velocity difference between the two features may be due to stripped gas leaving a ``hole'' in the distribution, making the secondary feature appear at a higher density than its immediate surroundings (we do not necessarily expect that the column density should smoothly vary over the entire cloud). Henceforth, we adopt the hypothesis that Pegasus is being ram pressure stripped. Following \cite{gunn1972}, material will be ram pressure stripped from a galaxy if the density of the surrounding medium, $n_{IGM} \gtrsim \left(2\,\pi\,G\,\Sigma_{T}\,\Sigma_{HI}\right)/\left(\mu v^2\right)$. $\Sigma_{T}$ is the total surface density (stars plus gas), $\Sigma_{HI}$ is the column density of HI and $v$ is the relative velocity of the galaxy to the medium. Thus, \begin{eqnarray} n_{IGM} &\simeq& 3.7 \times 10^{-6}\,\rm{cm}^{-3} \left(\frac{\rm{100\,km\,s^{-1}}}{v}\right)^2 \nonumber\\ &&\left(\frac{\Sigma_{HI}}{0.1 \times 10^{20}\,\rm{cm}^{-2}}\right)^2~, \end{eqnarray} \noindent where we take the mean particle mass $\mu = 0.75\,m_p$ for fully ionized media. We approximate the Local Group space velocity of Pegasus as $v \sim \sqrt{3}\,\sigma_{LG} \sim 100$\,km\,s$^{-1}$ where $\sigma_{LG} \sim 60$\,km\,s$^{-1}$ is the Local Group line-of-sight velocity dispersion (\citealt{sandage1986}). HI at a column density much lower than $\Sigma_{HI} \sim 0.1 \times 10^{20}$\,cm$^{-2}$ has been stripped from Pegasus, implying that this is a reasonable lower limit for use in this calculation. We adopt $\Sigma_T = \Sigma_{HI} \left(1 + M_\star/M_{HI}\right)$, where $M_\star \sim 1.24 \times 10^7\,M_\odot$ is the stellar mass of Pegasus (Table~1). This seems reasonable; the surface brightness of Pegasus is $25$\,mags\,arcsec$^{-2}$ at a radius of $r = 1.5^\prime$ on the minor axis (\citealt{nilson1973,devaucouleurs1991}), corresponding to a stellar surface density of $\Sigma_\star \sim 4 \times 10^{20}$\,cm$^{-2}$. This is approximately equivalent to the stellar-to-gas mass ratio multiplied by the gas surface density ($M_\star/M_{HI} \times \Sigma_{HI}$) at $r = 1.5^\prime$. These values yield $n_{IGM} \sim 3.7 \times 10^{-6}$\,cm$^{-3}$. However, given the uncertainties involved, it is entirely plausible that the value of $n_{IGM}$ could be at least an order of magnitude larger than in Equation~1. \subsection{Consequences} What is the source of the material that is stripping Pegasus? The bottom right panel of Figure~1 shows the distances of Pegasus to its nearest gas-rich neighbours. The dwarf neighbours are unlikely to be the source of the stripping medium; not only is the required mass of gas unrealistically large (an ejected spherical shell $\sim 1$\,kpc thick with a radius of $\sim 300$\,kpc would have a mass $> 3 \times 10^6\,M_\odot$ at a density of $n_{IGM}$) but the energy required is too large for a dwarf galaxy to reasonably provide. Alternatively, the gas could be associated with M31. From observations of the Magellanic stream, \cite{murali2000} estimate that the density of the Milky Way halo at the stream must be $\lesssim 10^{-5}$cm$^{-3}$, although \cite{stanimirovic2002} estimate $\sim 10^{-4}$ cm$^{-3}$. If the gas density in the halo of M31 is similar, then not only would M31 need to have a very extended corona, but its density would need to decrease very slowly with radius. Indeed, if the Milky Way has a similarly extended corona, then the two will overlap and the result may be observationally indistinguishable from a Local Group medium. The isolation of Pegasus raises the strong possibility that the stripping medium is associated with the Local Group, rather than individual galaxies within the group. Clusters of galaxies have such media, and observations of Holmberg~II imply the presence of an intra-group medium in the M81 group (\citealt{bureau2002}). The density of the intra-group medium implied in Equation~1 is of the same order as the density of the medium responsible for local OVI absorption detected by \cite{nicastro2002,nicastro2003} and \cite{sembach2003}, which they suggest is associated with either a Milky Way corona or a Local Group medium. Our result favors the latter interpretation. Theoretically, $\sim 30\,\%$ of baryons in the Local Volume are expected to be in a warm/hot phase ($T \sim 10^5 - 10^6\,$K; \citealt{kravtsov2002}); this is likely concentrated around galaxies and galaxy groups as an intra-group medium. If the stripping medium pervades the Local Group, why do more dwarf galaxies not show evidence of ram pressure stripping? \cite{lin1983} suggest that all the dSphs have been stripped in this fashion, (although \cite{mayer2006} show that ram-pressure stripping by itself is insufficient to remove all the gas from a dIrr). It is possible that the Local Group medium will be clumpy and perhaps Pegasus is passing through a region of higher density compared to other dIrrs. Alternatively, Pegasus could be falling into and interacting with the Local Group for the first time, as has recently been speculated for two dSph galaxies at large radii from M31 (And~XII, \citealt{chapman2007}; AndXIV, \citealt{majewski2007}). However, the reason why only Pegasus currently shows signs of ram-pressure stripping is unlikely to be known until such time as the masses and orbits of the dIrrs have been determined. Given the distances of these galaxies, this will be some time yet. | 7 | 10 | 0710.4582 |
0710 | 0710.3477_arXiv.txt | We use Monte-Carlo simulations, combined with homogeneously determined age and mass distributions based on multi-wavelength photometry, to constrain the cluster formation history and the rate of bound cluster disruption in the Large Magellanic Cloud (LMC) star cluster system. We evolve synthetic star cluster systems formed with a power-law initial cluster mass function (ICMF) of spectral index $\alpha =-2$ assuming different {\dts}s. For each of these cluster disruption time-scales we derive the corresponding cluster formation rate (CFR) required to reproduce the observed cluster age distribution. We then compare, in a ``Poissonian'' $\chi^2$ sense, model mass distributions and model two-dimensional distributions in log(mass) vs. log(age) space of the detected surviving clusters to the observations. Because of the bright detection limit ($M_V^{\rm lim} \simeq -4.7$ mag) above which the observed cluster sample is complete, one cannot constrain the characteristic cluster disruption time-scale for a $10^4$ M$_\odot$ cluster, $t_4^{\rm dis}$ (where the disruption time-scale depends on cluster mass as $t_{\rm dis} = t_4^{\rm dis} (M_{\rm cl} / 10^4 {\rm M}_\odot )^\gamma$, with $\gamma \simeq 0.62$), to better than a lower limit, $t_4^{\rm dis} \ge 1$\,Gyr. \\ We conclude that the CFR has been increasing steadily from 0.3 clusters Myr$^{-1}$ 5 Gyr ago, to a present rate of $(20-30)$ clusters Myr$^{-1}$, for clusters spanning a mass range of $\sim 100-10^7$ M$_\odot$. For older ages the derived CFR depends sensitively on our assumption of the underlying CMF shape. If we assume a universal Gaussian ICMF, then the CFR has increased steadily over a Hubble time from $\sim 1$ cluster Gyr$^{-1}$ 15 Gyr ago to its present value. On the other hand, if the ICMF has always been a power law with a slope close to $\alpha=-2$, the CFR exhibits a minimum some 5 Gyr ago, which we tentatively identify with the well-known age gap in the LMC's cluster age distribution. | \label{sec:intro} The mass and age distributions of star cluster systems contain the (fossil) records of their formation conditions. They are therefore among the best tracers of the star-formation histories of their host galaxies available to observers. It is important to realise, however, that one needs to understand both the dominant internal and external evolutionary processes in order to disentangle this formation record, and hence obtain a glimpse of the initial conditions required for star cluster formation. The effects of stellar evolution in a given star cluster (which can be approximated by a ``simple'' stellar population once the cluster has reached an age that is well in excess of its formation time-scale) are rather well understood, whereas we have only recently begun to make major {\it quantitative} inroads into understanding the environmental effects leading to star cluster ``weight loss'' (i.e., the preferential depletion of the low-mass component of the cluster's stellar mass function caused by tidal stripping and the ejection of stars owing to internal two-body relaxation) and -- eventually -- disruption. Estimates of the characteristic cluster disruption time-scales in various star cluster environments have been calculated by Boutloukos \& Lamers (2003), de Grijs et al. (2003a,b,c), Gieles et al. (2005) and de Grijs \& Anders (2006), among others (see also Lamers et al. 2005a,b). Specifically, Boutloukos \& Lamers (2003) and Lamers et al. (2005b) show that the cluster age distribution and cluster mass function approximate a double power law when a cluster system is affected by both fading and secular dynamical evolution. Knowledge of the detection limit in terms of cluster mass vs. age, combined with the estimate of the cluster age or mass at the ``break point'' of the integrated cluster age (mass) distribution (see, e.g., fig. 1 in Boutloukos \& Lamers 2003, where the ``break points'' are referred to as $t_{\rm cross}$ and $M_{\rm cross}$ in the age and mass distributions, respectively) then leads to the typical cluster disruption time-scale. Their analysis, however, explicitly builds on the assumption of a constant cluster formation rate (CFR, i.e. the number of clusters formed per linear time interval ${\rm d}N/{\rm d}t$ is constant in time) as a function of time. In that context, the (poorly-known) time-variable CFR of the LMC cluster system hampers such an analysis. If only the observed cluster age distribution is known, then we are left with a degeneracy between the CFR and the disruption time-scale, in the sense that one cannot distinguish between a low CFR combined with slow secular dynamical evolution on the one hand, and a vigourous CFR combined with cluster disruption occurring on a rapid time-scale on the other. The star cluster system in the Large Magellanic Cloud (LMC) has the potential of providing strong constraints to the theory of star cluster disruption as a function of environment, since it is composed of the largest resolved cluster system spanning {\it both} a reasonable mass range ($\sim 10^2 - 10^6$ M$_\odot$; cf. Hunter et al. 2003, hereafter H03; de Grijs \& Anders 2006, and references therein) {\it and} an age range from a few Myr to $\sim 13$ Gyr available. In addition, thanks to the LMC's proximity, we have been able to obtain observations -- and derived the age and mass distributions -- of a sufficiently large cluster sample to allow a statistical approach to its evolution (e.g., H03; de Grijs \& Anders 2006; see also Sect. 2). On the basis of the few star clusters systems analysed in detail to date, including M51 and the Antennae interacting system, Bastian et al. (2005), Fall et al. (2005) and Fall (2006) suggest that the early evolution of star cluster systems is most likely characterised by a rapid, largely mass-independent ``infant mortality'' phase, at least for masses $\ga 10^4$ M$_\odot$ (see also de Grijs \& Parmentier 2007; and references therein), combined with ``infant weight loss''(the loss of stars caused by rapid, early gas expulsion; cf. Weidner et al. 2007), the effects of which are enhanced by stellar evolutionary mass loss. In this scenario, this early phase, which ends when clusters reach an age of $\sim 40$ to 50 Myr (e.g., Goodwin \& Bastian 2006), would then be followed by (mass-dependent) secular evolution. The early, rapid cluster disruption process results from the expulsion of the intracluster gas due to adiabatic or explosive expansion driven by stellar winds or supernova activity (Mengel et al. 2005; Bastian \& Goodwin 2006; Goodwin \& Bastian 2006; see de Grijs \& Parmentier 2007 for a review). Star clusters are expected to settle back into virial equilibrium $\sim 40$ to 50 Myr after gas expulsion (Goodwin \& Bastian 2006). In our analysis in this paper we will therefore exclude star clusters younger than 50 Myr, since our main purpose is to derive the characteristic (mass-dependent) time-scale of cluster disruption in the LMC driven by secular evolution only. In a follow-up paper (Goodwin et al., in prep.), we will discuss the evolution of the LMC cluster system on the shortest time-scales relevant to the infant mortality and infant weight loss scenarios. In de Grijs \& Anders (2006), we found that the LMC's CFR has been roughly constant outside of the well-known age gap between $\sim 3$ and 13 Gyr, when the CFR was a factor of $\sim 5$ lower (assuming a roughly constant rate during this entire period). Based on this observation as our main underlying assumption, we used a simple approach to derive the characteristic cluster disruption time-scale in the LMC, for which we found that $\log(t_4^{\rm dis} {\rm yr}^{-1}) = 9.9 \pm 0.1$, where $t_{\rm dis} = t_4^{\rm dis} (M_{\rm cl}/10^4 {\rm M}_\odot)^{0.62}$ (Boutloukos \& Lamers 2003; Baumgardt \& Makino 2003; Gieles et al.~2005). We argued that this was consistent with earlier, preliminary work done on a smaller cluster sample: Boutloukos \& Lamers (2002) found $\log( t_4^{\rm dis} {\rm yr}^{-1} ) = 9.7 \pm 0.3$ for a smaller sample of 478 clusters within 5 kpc from the centre of the LMC, in the age range $7.8 \le \log(\mbox{age yr}^{-1}) \le 10.0$. We also considered our result qualitatively consistent with Hunter et al. (2003), who noticed very little destruction of clusters at the high-mass end. This long characteristic disruption time-scale would imply that hardly any of our LMC sample clusters are affected by significant disruptive processes, so that we are in fact observing the {\it initial} cluster mass function (CMF). However, a close inspection of fig. 6 of de Grijs \& Anders (2006) highlights an apparent contradiction. The ``crossing time'', $t_{\rm cross}$, defined by the crossing point between the best-fitting lines describing the number of clusters per unit time-scale that are mostly affected by fading of their stellar populations and those that are undergoing significant secular disruption, seems to imply that a more appropriate time-scale for the disruption of the LMC cluster system may be of order $\log( t_4^{\rm dis} {\rm yr}^{-1} ) \simeq 8.9$. Since this implies a downward adjustment of the characteristic cluster disruption time-scale in the LMC by up to an order of magnitude, we decided to re-investigate the LMC's cluster disruption history. Here, we approach this problem from a different angle, by running a large number of Monte-Carlo simulations in which we vary the cluster disruption time-scale. Meanwhile, for each of these cluster disruption time-scales we derive the corresponding CFR required to reproduce the observed cluster age distribution. We then match the observed cluster mass distribution, integrated over time, and the observed two-dimensional distribution of the detected surviving clusters in the log(mass) vs. log(age) plane to the model results. $\chi^2$ fit estimates are used to quantify which cluster disruption time-scale and, therefore, which cluster formation history, best describes the presently available data. This paper is organised as follows. In Section \ref{sec:data} we justify our choices used in the data analysis leading to the cluster age and mass distributions used in the remainder of the paper. Section \ref{sec:synth} discusses our basic assumptions in constructing synthetic cluster populations, which we then use in Section \ref{sec:disr} to explore the range of characteristic cluster disruption time-scales allowed by the data. In Section \ref{sec:comp} we highlight the importance of properly understanding the data's completeness limit, and use this in Section \ref{sec:agap} to constrain possible variations in the CFR over time. In Section \ref{sec:1vs10}, we assess precisely what we would need to fully and unambiguously constrain $t_4^{\rm dis}$. Our results and conclusions are summarized in Section \ref{sec:conc}. \begin{figure*} \begin{minipage}[t]{\linewidth} \centering\epsfig{figure=obs_age_mass.eps, width=\linewidth} \end{minipage} \caption{Distribution of the LMC star cluster sample of de Grijs \& Anders (2006) in the [$\log({\rm age}),\log(M_{\rm cl})$] plane. The filled triangles correspond to the vertical dashed lines in the individual panels of Fig.~\ref{fig:compl} (upright triangles: left-hand column; upside-down triangles: right-hand column). For subsequent cluster age ranges (in steps of 0.25 dex and 0.5 dex wide) they trace the mass limit below which the sample becomes incomplete (see section \ref{sec:synth} for details). They are therefore considered tracers of the fiducial detection limit (thick dash-dotted line with squares), which corresponds to $M_V^{\rm lim}=-4.7$ mag (based on the {\sc galev} mass-to-light ratios for ``simple'' stellar populations). The three thin dash-dotted lines, labelled `[1]', `[2]' and `[3]', are the detection limits corresponding to $M_V^{\rm lim}=-5.2$, $M_V^{\rm lim}=-4.2$ and $M_V^{\rm lim}=-3.5$ mag, respectively. They are therefore equivalent to the thick dash-dotted line shifted vertically by, respectively, $\Delta \log (M_{\rm cl})=0.2, -0.2$, and $-0.48$. The lower dash-dotted curve (`[3]') is the $M_V^{\rm lim}=-3.5$\,mag fading limit of H03. The thick dashed lines represent the cluster disruption limits for $\log (t_4^{\rm dis} {\rm yr}^{-1})=8.1, 9.0$ and $9.9$ (labelled `[a]', `[b]' and `[c]', respectively). The age range on which we focus in this paper is bracketed by vertical solid lines, at $\log (\mbox{age yr}^{-1})=7.7$ and $\log (\mbox{age yr}^{-1})=9.2$; clusters brighter than $M_V^{\rm lim}=-4.7$ mag in that age range are represented by crosses} \label{fig:obs_am} \end{figure*} \begin{figure*} \begin{minipage}[t]{\linewidth} \centering\epsfig{figure=MF_agebin_completeness.eps, width=\linewidth} \end{minipage} \caption{Observed cluster mass functions for the age ranges included at the top of each panel. In each panel, the vertical dashed line is the mass limit bracketing 25 and 75 per cent of the cluster subsample on either side. This is a good proxy to the cluster mass at the turn-over of each CMF. For each age range, turn-over masses are indicated in ($\log (\rm age),\log (M_{\rm cl})$) space as filled triangles in Fig.~\ref{fig:obs_am}. The mass limits defined by the dashed lines evolve with time following a line of constant luminosity, at $M_V=-4.7$ mag. This implies that the decrease in cluster numbers observed for each CMF below its turn-over mass (i.e. below the vertical dashed line) is mostly driven by incompleteness effects} \label{fig:compl} \end{figure*} | \label{sec:conc} In this paper, we have carried out numerous detailed Monte-Carlo simulations aimed at constraining the cluster formation history and the rate of bound cluster disruption in the LMC star cluster system. We considered only clusters older than 50\,Myr in order not to determine erroneously short cluster disruption time-scale as a result of the inclusion of the infant mortality and infant weight loss evolutionary phases. Our data are based on the $UBVR$ photometry of H03 (obtained from Massey's (2002) survey of the Magellanic Clouds), for which de Grijs \& Anders (2006) obtained homogeneously determined age and mass estimates. We estimated a fiducial detection limit above which the cluster sample is (fairly) complete from the CMF as a function of age, at $M_V^{\rm lim} = -4.7 {\rm mag}$. This is significantly brighter than H03's fading limit. We consider only clusters brighter than this limit, in order to avoid severe statistical incompleteness effects. We have evolved synthetic star cluster systems characterised by constant CFRs assuming 20 different {\dts}s. The CFR was adjusted to reproduce the observed age distribution. By doing so, we are likely to loose sensitivity to CFR variations occurring on a time-scale shorter than the width of the age range corresponding to each logarithmic age bin (i.e., more likely at old age where our age bins span greater linear age ranges than at young age). However, the general behaviour of the CFR, as well as the CFR averaged over age are robustly recovered (see Section \ref{sec:1vs10}). We then compare, in a ``Poissonian'' $\chi^2$ sense the modelled mass distribution and the modelled [$\log({\rm age}),\log(M_{\rm cl})$] distribution to the observations. We show that because of the bright detection limit at $M_V^{\rm lim}= -4.7$ mag, one cannot constrain $t_4^{\rm dis}$ to better than a lower limit, $t_4^{\rm dis} \ge 1$\,Gyr. The tightest constraints are set by the CMF integrated over age. The $\chi^2$ test applied to the distribution of points in [$\log({\rm age}),\log(M_{\rm cl})$] space turns out to be a poor diagnostic tool. This is probably related to the low density of points in each cell of the [$\log({\rm age}),\log(M_{\rm cl})$], compared to the density of points in each bin of the integrated CMF, leading to smaller Poissonnian error bars and therefore a better constrained model in the latter case. Our range of {\dts} estimates is robust with respect to model variations, such as of the upper limit of the initial cluster mass range, the location of the grid in [$\log({\rm age}),\log(M_{\rm cl})$], and the size of cells in this grid. We have shown that should the detection limit be underestimated, artificially shortened {\dts}s would result. This is so because there is a degeneracy between incompleteness and secular evolution, i.e., the fading-driven turn-over in a CMF in a given age bin is interpreted as resulting from secular evolution, leading to a shortened cluster disruption time-scale. Having set the best possible constraints on $t_4^{\rm dis}$, we explored the corresponding CFR, in particular considering $t_4^{\rm dis}=1$ and 10 Gyr. The CFR has been increasing steadily from about 0.3 clusters Myr$^{-1}$ 5 Gyr ago, to a present rate of $(20-30)$ clusters Myr$^{-1}$, for clusters spanning an initial mass range of $\sim 100-10^7$ M$_\odot$. The CFRs inferred for both disruption time-scales differ by at most a factor of three (Fig.~\ref{fig:ad_all}). At older ages however, the situation becomes unclear. The uncertainty in the CFR as a result of the uncertainty on $t_4^{\rm dis}$ increases. In addition, the overall temporal behaviour of the CFR depends on the shape of the ICMF of the oldest, globular cluster-like objects. If this is the universal Gaussian ICMF, then the CFR has increased steadily over a Hubble time from $\sim 1$ cluster Gyr$^{-1}$ 15 Gyr ago to its present value. On the contrary, if the ICMF has always been a power law with a slope close to $\alpha=-2$, the CFR exhibits a minimum some 5 Gyr ago. Our results may be related to the orbital history and dynamics of the LMC with respect to both the SMC and the Milky Way, although this remains poorly constrained because of a lack of proper motion measurements with the required accuracy (Besla et al.~2007). Additionally, we note that interactions between the Clouds and between the Clouds and the Galaxy, while affecting their star-formation history, also affect the {\it structure} of the Magellanic Clouds (Bekki \& Chiba 2005). This may have induced temporal variations in the {\dts} over the past Hubble time, rendering the {\dts} estimate at old age even more uncertain. Finally, we have investigated which strategy should be adopted in the future in order to better constrain the {\dts} in the LMC. Specifically, we have generated synthetic cluster populations defined by a given cluster formation history, ICMF and various dissolution rates. After the addition of Gaussian noise to mimic an observational situation, we processed these simulated data sets in the same way as the actual LMC data. We demonstrate that {\it if} the {\dts} is known, then the CFR can be derived accurately. We confirm our inability to distinguish $t_4^{\rm dis}=1$ Gyr from 10 Gyr because of the bright detection limit. With such a bright detection limit, the expected turn-over in the CMF caused by dynamical evolution is not detected, for any cluster age range. As a result, only a lower limit to the {\dts} can be retrieved, i.e., we can exclude all $t_4^{\rm dis}$ that are short enough to lead to a turn-over above the detection limit. To obtain age and mass estimates for an LMC star cluster sample complete above $M_V^{\rm lim}=-3.5$ is desirable to more easily distinguish between $t_4^{\rm dis}=1$ Gyr and $t_4^{\rm dis}=10$ Gyr. | 7 | 10 | 0710.3477 |
0710 | 0710.3180.txt | We present new optical spectroscopy and photometry, 2MASS infrared observations and 24 years of combined AAVSO and AFOEV photometry of the symbiotic star candidate \ae. The long-term light curve is characterized by outbursts lasting several years and having a slow decline of $\sim 2 \times 10^{-4}$ mag/day. The whole range of variability of the star in the $V$ band is about 4 magnitudes. The periodogram of the photometric data reveals strong signals at $\sim$ 342 and 171 days. The presence of the emission feature at $\lambda$ 6830 \AA~ at minimum and the detection of absorption lines of a $\sim$ K5 type star confirm the symbiotic classification and suggest that AE\,Cir is a new member of the small group of s-type yellow symbiotic stars. We estimate a distance of 9.4 kpc. Our spectrum taken at the high state shows a much flatter spectral energy distribution, the disappearance of the $\lambda$ 6830 \AA~ emission feature and the weakness of the He\,II 4686 emission relative to the Balmer emission lines. Our observations indicate the presence of emission line flickering in time scales of minutes in 2001. The peculiar character of \ae~ is revealed in the visibility of the secondary star at the high and low state, the light curve resembling a dwarf nova superoutburst and the relatively short low states. The data are hard to reconciliate with standard models for symbiotic star outbursts. %Another peculiarity is the detection of features of the cool star at the low and high states. %We discuss the multi-wavelength data and suggest that the ellipsoidal-like variability is due to %fluorescence of the cool star atmosphere that is irradiated by high energy photons from the hot component. %We observe a 40-day length eclipse-like episode in %our photometry, that we interpret as self-occultation of the irradiated hemisphere of the red star. Then, assuming a white dwarf %hot component of 1 M$\odot$ we find that the red giant probably fills their Roche lobe and has a radius $R_{g} \approx$ 100 R$_{\odot}$. %From independent methods we have found a distance to AE\,Cir of 14 $\pm$ 3 kpc. %Our observations suggest that accretion onto a compact object %in a semi-detached binary could play a major role in the high states and short-term variability observed in AE\,Cir. | AE\,Cir (S32 = StHA32, $\alpha_{2000}$ = 14:44:52.0, $\delta_{2000}$ = -69:23:35.9) was classified as a RCB star in the Fourth General Catalog of Variable Stars. This classification was rejected by Kilkenny (1989), who found strong emission lines of H\,I and weaker emission lines of He\,II and other species in one spectrogram spanning a wavelength range of 3400-5100 \AA. Kilkenny noted the lack of forbidden lines and the anomalous strong He\,II 4686 emission for this object, and suggested a symbiotic classification, although no lines of the cool component were detected. He also mentions the photometric variability of the star, ranging from 12 to 14 mag. (from visual estimates) in hundreds of days. $B,V$ photoelectric measures by Lawson \& Cottrell (1990) taken in an interval of 107 days show the star with $V$ magnitude between 13.52 and 14.51 and color $B-V$ between 0.95 and 1.43, the object being redder when fainter. Based on the spectroscopy by Kilkenny, AE\,Cir was listed as a suspected symbiotic in the catalog of Belczynski et al. (2000). Symbiotic stars have been reviewed recently by Miko{\l}ajewska (2007). In this paper we present new photometry and low and high-resolution spectra of AE\,Cir and investigate their emission line properties and symbiotic nature. In Section 2 we give details of our spectroscopic observations and show the methods used in data reduction and analysis. In that section we also analyze available visual long-term photometric records along with our own CCD photometry. In Section 3 we analyze our data, giving the main results of our research. In Section 4 we discuss our results in the context of symbiotic stars and present our conclusions in Section 5. | In this paper we have investigated the nature of the symbiotic star candidate AE\,Cir. We have analyzed new optical photometric and spectroscopic data, 2MASS infrared photometry and 24 years of visual photometry in an integrated way in order to get a better understanding of the system. Our conclusions can be summarized as follows:\\ \begin{itemize} \item The symbiotic nature is confirmed. This result is based on the detection of the $\lambda$ 6830 \AA~ emission feature and the spectral signatures of a cool stellar component of spectral type $\sim$ K5. \item The spectral type $\sim$ K5 and the spectral and photometric properties indicate that \ae~ is a member of the small group of yellow symbiotic stars of the s-type. \item The light curve is characterized by outbursts lasting $\sim$ 4000 days and overall amplitude of variability about 4 magnitudes. The outbursts show a slow decline of $\sim 2 \times 10^{-4}$ mag/day. The duration of the low state is about 38\% the high state. \item A strong signal at 342 $\pm$ 15 days is detected in the light curve. The light curve folded with this period shows two broad minima with different amplitude. \item The spectrum in the low state is characteristic of a symbiotic star without forbidden lines and very strong He\,I\,4868 emission. At maximum the emission feature at $\lambda$ 6380 \AA~ disappears, the spectrum becomes flat and the relative intensity of the He\,I 4686 \AA~line, compared with the Balmer lines, becomes lower. At the same time the overall emission line spectrum shows smaller equivalent widths. Signatures of the cool companion are observed at the high and low state. \item We observe H\,I, He\,I and He\,II emission line variability in time scales of minutes in 2001 being larger in the line centers. The Na\,D and $\lambda$ 6820 \AA~ lines do not follow this variability. We also observe larger broadening and asymmetry in the Balmer lines compared with others emission lines in this epoch. \item The data suggest that full visibility of the red giant should occur at the low state at $V \sim 15.5$. For a typical K5 giant behind the galactic disc this would indicate a distance of 9.4 kpc. This should indicate that \ae~ is about 1.4 kpc below the galactic plane. \item We suggest that an eclipse-like event could be interpreted as self-occultation of the irradiated red giant hemisphere. In this case we found that it is possible that the red giant fills its Roche-lobe. Assuming a white dwarf hot component of 1 M$_{\odot}$, we estimate a radius of 99 R$_{\odot}$, a mass of 1.1 M$_{\odot}$ and M$_{V}$ = -2.4 for the red giant at the epoch of the SMARTS observations. \item The atypical character of \ae~ is revealed in the long time passing in outburst and the dwarf nova superoutburst shape of the light curve. Another atypical feature is the rather strong TiO absorption band around $\lambda$ 6300 \AA~ at the high state. This could indicate that the red giant and the hot component are partially obscured at minimum or the brightness of the secondary star is not constant. Obscuration hardly explains the large line emission observed at minimum and the long brightness decay after maximum whereas variations of the red giant brightness imply unrealistic changes in the stellar radius. \item At present, the available data for AE\,Cir are hard to reconciliate with canonical models for symbiotic outbursts. %\item {\bf The similar visibility of the secondary star at the low and high state probably rules out the %hypothesis of the hot component outbursts and suggests an scenario where the %outbursts are due to thermal instabilities in an accretion disc fed by a Roche lobe filling giant star. %In this case the low state spectrum is dominated by the giant star and emission lines arising from wind accretion. %On the contrary, the hot state is characterized by a blue continuum arising from the hot and luminous disc in outburst %and a diminishing of the spectral features arising from the accreted wind. Irradiation of the secondary star by the %outbursting disc, and eventually stellar pulsation, could explain its variable brightness. The disappearance of the $\lambda$ %6830 \AA~ feature at the high state could be explained if the O\,VI photon forming region is occulted by the accretion disc %during outburst.} \end{itemize} | 7 | 10 | 0710.3180 |
0710 | 0710.2317_arXiv.txt | The goal of this research is to investigate how well various turbulence models can describe physical properties of the upper convective boundary layer of the Sun. An accurate modeling of the turbulence motions is necessary for understanding the excitation mechanisms of solar oscillation modes. We have carried out realistic numerical simulations using several different physical Large Eddy Simulation (LES) models (Hyperviscosity approach, Smagorinsky, and dynamic models) to investigate how the differences in turbulence modeling affect the damping and excitation of the oscillations and their spectral properties and compare with observations. We have first calculated the oscillation power spectra of radial and non-radial modes supported by the computational box with the different turbulence models. Then we have calculated the work integral input to the modes to estimate the influence of the turbulence model on the depth and strength of the oscillation sources. We have compared these results with previous studies and with the observed properties of solar oscillations. We find that the dynamic turbulence model provides the best agreement with the helioseismic observations. | Dominant acoustic sources within the Sun are generated by strong fluctuations in the outer convective layers. Turbulent motions stochastically excite the resonant modes via Reynolds stresses and entropy fluctuations. The dominant driving comes from the interaction of the nonadiabatic, incoherent pressure fluctuations with the coherent mode displacement \citep{Nordlund2001}. The modes excitation sources occur close to the surface, mainly in the intergranular lanes and near the boundaries of granules \citep{Stein2001}. Thus an accurate modeling of the turbulence motions is necessary to understand the excitation mechanisms of solar oscillation modes. The correct choice of turbulence model is also important in many other astrophysical simulations. The objective of this research is to study the influence of turbulence models on the excitation mechanisms by means of realistic numerical simulations. We have compared different physical Large Eddy Simulation (LES) models (Hyperviscosity approach, Smagorinsky, and dynamic models) to show the influence on the damping and excitation of the oscillations. The organization of this paper is as follows. In \S2, we describe the main lines of the code and the different turbulence models. The kinetic energy of the radial modes obtained with the different turbulence models are presented in \S 3. Then a comparison of the results obtained with the different turbulence models for the non-radial modes is given in \S 4. The work integral input to the modes is calculated in \S5 in order to estimate the influence of the turbulence models on the depth of the oscillation sources. | The goals of this research was to investigate how well various turbulence models can describe the convective properties of the upper boundary layer of the Sun and to study the excitation and damping of acoustic oscillations. Results obtained with the hyperviscosity approach have been compared with those obtained with the Smagorinsky and dynamic turbulence models. We have seen that the dissipation is very high with the Smagorinsky model while the hyperviscosity approach and dynamic modes give similar results. Besides we find that the dynamic turbulence model provides the best agreement with observations. | 7 | 10 | 0710.2317 |
0710 | 0710.5701_arXiv.txt | {} {We present the period analysis of unfiltered photometric observations of V5116~Sgr (Nova Sgr 2005 \#2) and we search for superhump candidates in novae remnants.} {The PDM method for period analysis is used. The masses of the novae componets are estimated from the secondary mass -- orbital period and primary mass -- decline time relations.} {We found that 13 nights of V5116~Sgr observations in the year 2006 are modulated with a period of $0.1238 \pm 0.0001$~d ($2.9712 \pm 0.0024$~h). Following the shape of the phased light curves and no apparent change in the value of the periodicity in different subsamples of the data, we interpret the period as orbital in nature. The binary system then falls within the period gap of the orbital period distribution of cataclysmic variables. From the maximum magnitude -- rate of decline relation, we estimate the maximum absolute visual magnitude of $M_{\rm Vmax} = -8.85 \pm 0.04$~mag using the measured value of decline $t_{\rm 2} = 6.5 \pm 1.0$~d. The mass-period relation for cataclysmic variables yields a secondary mass estimate of about $0.26 \pm 0.05~{\rm M}_{\rm \odot}$. We propose that V5116~Sgr is a high inclination system showing an irradiation effect of the secondary star. No fully developed accretion disc up to the tidal radius with the value lower than $3.5~10^{10}$~cm is probable. The mass ratio was estimated in a few novae and the presence or absence of superhumps in these systems was compared with the mass ratio limit for superhumps of about 0.35. We found that in the majority of novae with expected superhumps, this variability has not been found yet. Therefore, more observations of these systems is encouraged.} {} | Novae are a subclass of cataclysmic variable stars. In these interracting binaries, the white dwarf is accreting the matter transfered from the secondary star. The accretion disc may be formed in the non magnetic case. The intermediate polar systems have a truncated disc and polar systems have magnetic field strong enough to prevent the disc formation (see Warner 1995 for review). The accumulation of critical amount of accreted material onto the white dwarf surface results in a nova explosion. The distribution of orbital periods in cataclysmic variables shows a period gap between about 2 and 3 hours. Novae do not show this lack of objects in the mentioned interval. V5116~Sgr (Nova Sgr 2005 \#2) was discovered by Liller (2005) on 2005 July 4.049 UT. The nova had a magnitude $\sim 8.0$ on two red photographs. An unfiltered CCD image from 2005 July 5.085 UT showed the object at mag 7.2. The spectrum from 2005 July 5.099 UT showed H$_{\rm \alpha}$ with the FWHM of $\sim 970$ km~s$^{-1}$. The expansion velocity derived from the sharp P~Cyg profile was $\sim 1300$ km~s$^{-1}$. The position of the nova was measured by Gilmore and Kilmartin (2005) and Jacques (2005). Gilmore and Kilmartin (2005) searched for the nova precursor, but no convincing candidate has been found. Russell et al. (2005) performed a 0.8 -- 2.5 $\mu$m spectroscopy of the nova on 2005 July 15. The object showed emission lines of H~I, He~I, C~I, N~I, Ca~II and O~I with a FWHM $\sim$ 2200 km~s$^{-1}$. He~I showed P~Cyg profile at 1.0830 and 2.0581 $\mu$m. No thermal dust emission was observed. After the nova explosion the accretions disc is destroyed. The new disc is reformed by the stream of matter flowing from the secondary, interracting with itself and forming a ring in the circularisation radius. The disc starts to form by the viscous shearing. Matter losing angular momentum is moving invard and the excess of angular momentum is transported by the matter flowing outward. The invard moving matter touchs the white dwarf and the disc is reformed (Pringle 1981). If the white dwarf is magnetic, the matter interracts with the magnetosphere and is then conducted by the magnetic field to the magnetic poles. The interraction of the flow with the poles of the rotating star is observed as periodic signal with the spin period of the white dwarf. In the case of intermediate polars the spin period is usually much shorter than the orbital period (Patterson 1994, Hellier 1996). The polar systems without a disc are synchronous rotators, hence the spin period is equal or close to the orbital period (see e.g. Schmidt and Stockman 1991). Nova V1500 Cyg (polar system without a disc) changed the period from 0.141~d to 0.138~d and then stabilised at 0.140~d (Patterson 1978, 1979). The difference of 1.8\% between the rotation period of the white dwarf and the orbital period is ascribed to the effects of the nova explosion in 1975. The synchronisation can be then corrupted by the nova explosion, but the spin period remain still very close to the orbital period. This is in contrast with intermediate polars with disc presence. Searching for periods in novae allows to study the orbital distrubution and evolution of these systems. Currently, there are about 50 novae with known orbital periods (Warner 2002) with typical values ranging from 2 to 9 hours. The existence of the accretion disc or its renovation after the nova explosion is confirmed by the detection of the superhump period (see e.g. Retter et al. 1997, Kang et al. 2006a) or the spin period of a magnetic white dwarf in the case of intermediate polars (see e.g. Retter et al. 1998). Superhumps are caused by a precessing accretion disc generally in systems with mass ratio $M_{\rm 2}/M_{\rm 1} < 0.35 \pm 0.02$ (see Patterson et al. 2005 for review), in which the disc radius reaches the 3:1 resonance. The mass ratio indicates that systems with massive primaries and low mass secondaries are probable superhumpers. Novae in general possess high mass white dwarfs (see e.g. Warner 1995, Smith and Dhillon 1998, Webbink 1990) and short orbital periods sugest low mass secondaries. It is therefore meaningfull to search for superhump variability in novae with short orbital period. We have an ongoing program to observe novae with small telescopes to search for periodicities in their optical light curves. In this paper we report the detection of one periodicity in our photometric data ($P = 0.1238 \pm 0.0001$~d) of V5116~Sgr and we discuss the presence of superhumps in nova remnants. In Section~\ref{obs} we present our observational material. The long-term light curve with the period analysis of the data is presented in Sec.~\ref{data_anal} and in Sec.~\ref{disc} we discuss the long-term light curve behaviour ({Sec.~\ref{disc_v5116_l}}), the results of the period analysis (Sec.~\ref{disc_v5116_s}) and superhump search in nova remnants (Sec.~\ref{disc_SH}). | \label{disc} \subsection{Long-term variations of V5116~Sgr} \label{disc_v5116_l} The long-term light curve of the nova V5116~Sgr (Fig.~\ref{curve} -- bottom panel) shows a transition from smooth decline to probable oscillations. Several models has been suggested for such ``transient'' phase. One of them explains the transition as the time when the accretion disc is re-established (Retter 2002). The author propose a connection between this phase and intermediate polars. There is no indication of magnetic nature of the white dwarf in V5116~Sgr yet. Spin period of the white dwarf indicating the intermediate polar type of this ``transient'' phase in novae was detected for example in V4745~Sgr by optical observations (Dobrotka et al. 2006a), V1494~Aql (Drake et al. 2003) and V4743~Sgr by X-ray observations (Ness et al. 2003). If this ``transient'' phase interpretation of Retter (2002) is applicable in the case of V5116~Sgr, the accretion must be then re-established and the accretion disc should be present. \subsection{Short-term variations of V5116~Sgr} \label{disc_v5116_s} We have identified one periodicity in the light curve of the nova V5116~Sgr about 15 months after the maximum brightness. The value is $P = 0.1238 \pm 0.0001$~d ($2.9712 \pm 0.0024$~h). The upper limits of the period difference between two subset of data with $\simeq 16$ days of mean distance analysed in Fig.~\ref{power2} is $1.5~10^{-4}$~d which gives $|\dot{P}| \simeq 0.94~10^{-5}$. Following Patterson et al. (1993) the mean variation of the superhump period in SU~UMa superhumping sytems is $|\dot{P}| \simeq 3-10~10^{-5}$ (see their Table.~1). For the recurrent nova VY~Aqr the authors derived a variation of $|\dot{P}| \simeq 8.2~10^{-5}$. Our value of $|\dot{P}|$ comes probably from the errors of period measurements rather than from real period changes. The period seems then to be constant during our observations and the shape of the folded light curve suggests a primary and a secondary eclipse. We therefore propose that this periodicity is the orbital period of the binary system. Such a period is at the lower end of the mostly populated region of orbital periods in novae (Warner 2002). The dominant first harmonic frequency in the power spectra is a result of the clear structure primary -- secondary eclipse in the folded light curve which suggests a high inclination angle of the binary system. Using the orbital period of the system and equation (9) from Smith and Dhillon (1998) we obtain a rough estimate for the secondary star mass of $0.26 \pm 0.05$~M$_{\rm \odot}$. Using a mean white dwarf mass of $0.85 \pm 0.05$~M$_{\rm \odot}$ from Smith and Dhillon (1998), we find a mass ratio (secondary/primary mass) of $M_{\rm 2}/M_{\rm 1} = 0.3 \pm 0.1$. After the nova explosion, the hot white dwarf may heat and irradiate the cooler companion. The observed orbital light curve of the nova can be the result of the aspect variations or eclipses of the secondary due to heating from the hot primary and the asymmetry in the pulse profiles could be produced once the shape of the secondary is of a tear drop model. The irradiation effects in classical novae can also be detected long after the outburst stage (e.g., V1500~Cyg; Sommers and Naylor 1999, DN~Gem; Retter et al. 1999). Two eclipse like features are present in the folded light curve of V5116~Sgr (Fig.~\ref{folded}). The shape is similar to the light curve of V2540~Oph in 2003 (Ak et al. 2005) with a dip at phase $\sim 0.5$ ($\sim 33\%$ amplitude of the primary minimum for V2540~Oph and $\sim 70\%$ for V5116~Sgr). The authors concluded that V2540~Oph is likely a high inclination system either showing an irradiation effect or having a spiral structure in its accretion disc. Woudt and Warner (2003a) noted that one of the following requirements must be fulfilled for the large amplitude orbital modulations to be seen in the light curve of a recent nova in which the accretion disc does not dominate the luminosity of the system: 1) the disc is foreshortened but the secondary is seen (a high inclination angle), 2) the disc has small dimensions, 3) no disc (case of polars). In the case of V5116~Sgr the polar interpretation is possible because of the orbital period distribution of polars which peaks below 5 hours (Warner 1995), but nothing else indicate this option. The small dimension of the disc is supported by; 1) the short orbital period, 2) the post nova stage when the disc is reforming after explosion. A comparison to two novae within the period gap IM~Nor and DD~Cir with present irradiation effect (Woudt and Warner 2003b,c) can be made. The light curve of V5116~Sgr is different from those of IM~Nor and DD~Cir but similar to V2540~Oph. IM~Nor and DD~Cir showed very small dips at phase 0.5 interpretted as partial eclipses of the irradiated secondary by the disc or matter stream. The light curve shape depends on the disc radius and on the inclination angle. The deep secondary eclipses in V5116~Sgr require a large disc or a high inclination. The strength of the irradiation effect (white dwarf post-outburst temperature) can also play a significant role. The differences in the phase of the maxima in Fig.~\ref{folded}b,c can be due to spiral structures in the disc as mentioned before in regard with the nova V2540~Oph. The reconstruction of an accretion disc after the nova eruption is indicated by the discovery of the white dwarf spin or by the superhump period. A probable spin period of the white dwarf was detected $\sim 1$ year after the outburst in V4745~Sgr (Dobrotka et al. 2006a), $\sim 2.75$ years in V4743~Sgr (Kang et al. 2006b) and $\sim 15$ months after the maximum in V1425~Aql (Retter et al. 1998). Several systems show superhumps as early as two and a half months after the outburst like V4633~Sgr (a spin period is another option in this case, Lipkin et al. 2001) or two years after the outburst like V1974~Cyg (Retter et al. 1997). In our V5116~Sgr light curves extending $\sim 470$ days after outburst we did not find any photometric variations consistent with white dwarf spin modulation or superhump properties. We can not say anything about the magnetic nature of the white dwarf, but the components mass ratio indicates that the superhump existence is probable. The mass ratio using the primary mass average is not decisive. Using the derived secondary mass 0.26~M$_{\rm \odot}$ (Section~\ref{disc_v5116_s}) the mass ratio is lower than 0.35 for primary mass higher than $0.74$~M$_{\rm \odot}$. Taking all known primary masses in cataclysmic variables (Ritter and Kolb 2003), 63\% have higher or equal mass than $0.74$~M$_{\rm \odot}$. This probability value is not enought to make sure that superhumps are expected inV5116~Sgr, but the search for such variability could be fruitful. However V5116~Sgr is a very fast nova. According to equation (13) from Livio (1992) and taking $t_{\rm 3} = 20.2 \pm 1.9$~d derived in this paper (Section~\ref{t2t3}) we obtain a mass estimate of the white dwarf of $1.04 \pm 0.02$~M$_{\rm \odot}$ and thus $M_{\rm 2}/M_{\rm 1} = 0.25 \pm 0.05$. The presence of superhumps is then expected if a disc is fully developed up to tidal radius. The superhump period is a few percent longer or shorter than the orbital period. The possibility that the periodicity found in this paper is a superhump is rejected following the stability discussion and the shape of the folded light curve (Fig.~\ref{folded}). The mean shape of superhumps is typically an asymetric sinusoid and our data show typical eclipse like features. The presence of the disc in cataclysmic variables depends on the mass loss from the secondary. The matterial supplied from the secondary within the period gap depends on the strength of the magnetic braking driving the secondary out of thermal equilibrium (mass loss time scale shorter than the thermal time scale). The stronger the braking, the wider the gap will be and the higher is the upper end of the period gap. If the magnetic braking is low enought, the mass loss time scale may never become shorter than the thermal time scale. In the case of novae another condition is important. The nova explosion heats the secondary which leads to an enhanced mass transfer. Therefore the complete absence of accretion discs in novae within the period gap is not expected. \subsection{Searching for superhumps in other nova remnants} \label{disc_SH} The orbital period 2.462~h of the mentionned IM~Nor (Woudt and Warner 2003b) yields a secondary mass estimate of 0.20~M$_{\rm \odot}$ following Smith and Dhillon (1998). The decay time $t_{\rm 3} \simeq 50$~d (Kato et al. 2002) yields a primary mass estimate of $\simeq 0.86$~M$_{\rm \odot}$ (Livio 1992) giving a mass ratio $\simeq 0.23$. Woudt and Warner (2003b) concluded that they did not observe the superhump modulation. This conclusion with the mass ratio safely lower than 0.35 indicate that the existence of a fully developed disc is not probable. DD~Cir (Woudt and Warner 2003c) has a period of 2.339~h yielding a secondary mass estimate of 0.18~M$_{\rm \odot}$. The decay time $t_{\rm 2} \simeq 4.5$~d (Liller 1999, no other estimates are available) place the object in the class of very fast novae (similar to V5116~Sgr). The decay time $t_{\rm 3} \simeq 10$~d following the equation $t_{\rm 3} \simeq 2.75~t_{\rm 2}^{0.88}$ from Warner (1995) gives a rought white dwarf mass estimate of $\simeq 1.16$~M$_{\rm \odot}$ (mass ratio $< 0.2$). Woudt and Warner(2003b) interpretted the deep eclipse as obscuration of the disc and did not detect any periodicity near the orbital period in the period analysis. Therefore, following the tidal instability model it is again probable that the disc is not fully developed up to 3:1 resonance radius 3 years after the outburst. However the authors argued that the disc radius is 47 \% of the orbital separation. Using the third Kepler law we obtain a disc radius of $\sim 3.1~10^{10}$~cm. The disc in DD~Cir is then large enough to reach the 3:1 resonance radius calculated from the equation (3.39) from Warner (1995); $r_{\rm 3:1}\sim 3~10^{10}$~cm. The result is marginally in contrast with the absence of superhumps. Taking the derived values for V5116~Sgr we get $r_{\rm 3:1}\simeq 3.5~10^{10}$~cm as an upper limit of the disc radius taking the absence of superhumps into consideration. The boundaries of the eclipse in our light curve are not so clear as in the DD~Cir case, therefore not suitable to estimate the disc radius. Comparison to other systems follows the same way as in the case of IM~Nor and DD~Cir. We took components masses from the Catalogue of Ritter and Kolb (2003) and other masses are estimated from the orbital period (using Smith and Dhillon 1998) and $t_{\rm 3}$ time (using Livio 1992). We rejected systems with unknown or non measurable $t_{\rm 2}$ or $t_{\rm 3}$ time (peculiar light curve, strong oscilations in the decline) without catalogue mass values (no information on mass ratio, ex: RS~Car), magnetic novae (AM~Her systems without disc, ex: V1500~Cyg), systems with orbital period $> 10$ h (Smith and Dhillon fitting not adequate, systems with evolved secondary are then excluded too, ex: DI~Lac, V841~Oph, V723~Cas, GK~Per), systems with insuficient photometric observations to detailed period analysis (ex: V500~Aql, V446~Her, HZ~Pup, DY~Pup, CT~Ser). The final list of systems is in Table~\ref{systems}. \begin{figure} \includegraphics[width=90mm]{q_p.eps} \caption{Selected nova remnants in the mass ratio ($q$) -- orbital period ($P_{\rm orb}$) plane. Open circle -- non detected superhumps, filled circle -- detected superhumps. Panel a -- mass ratio from combined sourcees (catalogue and estimates from $P_{\rm orb}$, $t_{\rm 3}$ time), panel b -- mass ratio only from catalogue values (Ritter and Kolb 2003). The dashed line is the 0.35 limit.} \label{q_p} \end{figure} The results are depicted in Fig.~\ref{q_p}. The lower panel only shows the systems with known catalogue mass values and the upper panel is after the addition of estimated mass values in this work. Filled circles are systems with detected superhumps. The critical mass ratio of 0.35 is also shown. The results presented in the lower panel are as expected from the mass ratio limit for superhumps, but the statistical set is small. Adding other mass estimates changes the situation. 7 systems below or close to the limit 0.35 show superhumps but 14 systems do not. The estimated mass ratio has a mean difference of 0.2 from the catalogue values. The differences are distributed randomly. Therefore there is no reason to suspect that all mass ratios are systematically underestimated. The 7 systems with detected superhumps are safely located in comparing to the 0.35 critical value as expected by theory and are only a third of all systems falling below or close to the mass ratio of 0.35. The majority of all systems occupy the orbital period interval of 2 -- 4 hours. Following Patterson (2005) this interval has a $\sim$ 40 -- 90 \% probability to observe superhumps. 7 systems versus 14 from our investigation yields only $\sim 30$ \%. Superhumps are observed during the outbursts of short orbital period dwarf novae (SU~UMa stars) where the disc reaches the 3:1 resonance radius (superoutburst case). The disc during this active stage (outburst -- active and hot stage, quiescence -- not active and cold stage) reaches larger radius that in quiescence (see Lasota 2001 for review). The irradiation of the secondary by the central white dwarf in nova remnants can produce enhanced mass transfer rate strong enough to keep the disc in this stable hot active stage (as in the nova like stars) with the larger diameter (Hameury and Lasota 2005). In addition, superoutbursts in SU~UMa stars, when superhumps are observed, can be explained by the irradiation and the enhanced mass transfer (Smak 1995, 2004, Hameury et al. 2000, Schreiber et al. 2004). Therefore, permanent superhumpers are possible after a nova outburst. A few possible interpretation of the superhump lack can be: 1) not systematically applicable tidal instability theory; tidal torques are probably not the main and only condition of superhump existence as concluded by Hameury and Lasota (2005), 2) disc radius not developed up to tidal radius, 3) superhumps present but not detected, because of insufficient sets of data. | 7 | 10 | 0710.5701 |
0710 | 0710.5192_arXiv.txt | The group of 7 thermally emitting and radio-quiet isolated neutron stars (INSs) discovered by ROSAT constitutes a nearby population which locally appears to be as numerous as that of the classical radio pulsars. So far, attempts to enlarge this particular group of INSs finding more remote objects failed to confirm any candidate. We found in the 2XMMp catalogue a handful of sources with no catalogued counterparts and with X-ray spectra similar to those of the ROSAT discovered INSs, but seen at larger distances and thus undergoing higher interstellar absorptions. In order to rule out alternative identifications such as an AGN or a CV, we obtained deep ESO-VLT and SOAR optical imaging for the X-ray brightest candidates. We report here on the current status of our search and discuss the possible nature of our candidates. We focus particularly on the X-ray brightest source of our sample, 2XMM J104608.7-594306, observed serendipitously over more than four years by the XMM-Newton Observatory. A lower limit on the X-ray to optical flux ratio of $\sim$ 300 together with a stable flux and soft X-ray spectrum make it the most promising thermally emitting INS candidate. Beyond the finding of new members, our study aims at constraining the space density of this population at large distances and at determining whether their apparently high local density is an anomaly or not. | One of the outstanding results of ROSAT is the discovery of seven X-ray bright isolated neutron stars (INSs). These slowly rotating and radio-quiet neutron stars display thermal emission with $kT \sim$ 40 -- 110 eV, undergo little interstellar absorptions and are not associated with any supernova remnant (see reviews in \cite{tre00} and \cite{hab07}). Several have identified faint optical counterparts with $B \sim$ 25.8 -- 28.6. Proper motion studies (see Motch et al., these proceedings, and references therein) have shown that they are most probably young cooling neutron stars, with ages of a few 10$^5$ years. Their proximity and the apparent absence of strong non-thermal activity turn them into unique laboratories for testing radiative properties of neutron star surfaces in extreme conditions of gravitational and magnetic fields. Moreover, the possibility of measuring their distances through parallaxes \cite{ker07} or from the distribution of absorption on the line of sight \cite{pos07} can eventually bring important constraints on the debated equation of state of matter in neutron star interiors. In the solar neighborhood, ROSAT INSs are as numerous as young radio and $\gamma$-ray pulsars. It is not clear whether this group is homogeneous. In particular, the absence of radio emission can be either due to the presence of intense magnetic fields, indeed inferred from the measured spin down rates and cyclotron X-ray spectral features, or due to the fact that the radio pencil beam, which narrows at long spin periods, does not sweep over the earth. Considering that ROSAT had not enough sensitivity and spatial resolution to detect the thermal emission of distant sources, the population of cooling radio-quiet INSs could represent a considerable fraction of the total neutron star population of the Galaxy, undetectable in radio surveys \cite{mot07}. In any case, our knowledge of the overall population characteristics will remain highly unsatisfactory as long as only seven objects are known. \begin{table} \begin{tabular}{l c r r c r r} \hline \tablehead{1}{l}{t}{Source} & \tablehead{1}{c}{t}{Cand} & \tablehead{1}{r}{t}{RA} & \tablehead{1}{r}{t}{DEC} & \tablehead{1}{c}{t}{$R_{90}$} & \tablehead{1}{r}{t}{Count Rate} & \tablehead{1}{r}{t}{Mag}\\ & & (J2000) & (J2000) & arcsec & s$^{-1}$ & $R$\\ \hline 2XMM J104608.7-594306 & 065 & 10 46 08.7 & -59 43 06.1 & 1.33 & 0.060(4) & $>$ 25\\ 2XMM J121017.0-464609 & 164 & 12 10 17.1 & -46 46 11.2 & 2.60 & 0.027(6) & 20.3 \\ 2XMM J010642.3+005032 & 318 & 01 06 42.4 & +00 50 31.3 & 3.80 & 0.020(5) & 24.5 \\ 2XMM J214026.1-233222 & 364 & 21 40 26.2 & -23 32 22.3 & 1.90 & 0.0181(20) & 23.8 \\ 2XMM J125904.5-040503 & 604 & 12 59 04.6 & -04 05 02.3 & 1.90 & 0.0129(21) & 21.2 \\ 2XMM J125045.7-233349 & 681 & 12 50 45.7 & -23 33 47.7 & 3.30 & 0.0122(21) & 22.1 \\ \hline \end{tabular} \caption{INS candidates selected for optical investigation which have already been observed. Count rates are for the EPIC pn camera in the full XMM-Newton energy band (0.2--12.0 keV). We list the $R$ magnitude of the brightest object present in the error circles of the X-ray sources.\label{tab_cand}} \end{table} \begin{figure} \includegraphics[height=.28\textheight]{besteso_0010.eps} \caption{The positions of our INS candidates, selected from $\sim$ 7.5$\cdot$10$^4$ XMM sources, are shown in the HR$_1 \times$ HR$_2$ diagram (squares). The known ROSAT INSs (crosses) occupy the lowest (less absorbed) part of the diagram. Dashed lines denote soft absorbed blackbodies of different temperatures (50-200 eV) and column absorptions. The nine INS candidates selected for optical follow-up during the present year are highlighted with circles. The X-ray brightest and most promising candidate (source 65) can be noticed by its somewhat smaller error bars. \label{fig_bestcand}} \end{figure} | Our search for new thermally-emitting INSs in the 2XMMp catalogue has revealed a handful of interesting and previously unknown soft X-ray sources among which source 65 is, by far, the most promising candidate. The analysis of its X-ray emission, although based on archival data obtained with non-optimal configurations, reveals an intrinsically soft energy distribution, apparently stable on long time scales. The derived $N_{\textrm{H}}$ is consistent with that observed towards Eta Carinae and its cluster ($N_{\textrm{H}} \sim$ 3$\cdot$10$^{21}$ cm$^{-2}$). Scaling from RX J0720.4$-$3125 yields a distance of $\sim$ 3.9 kpc probably compatible with the 2.5 kpc assumed for the Carina Nebula \cite{2007ApJ...656..462S}. Optical follow-up observations failed to reveal counterparts brighter than $R \sim$ 25, supporting the idea that source 65 is a new thermally emitting INS. We have already obtained optical data for some of the other candidates. Preliminary analysis of these data reveals the presence of faint optical candidates within the error circles in all cases. We plan to use optical colour indices and spectra to identify these sources, thus characterizing this sample of soft objects strictly selected from more than 120 thousand entries present in the 2XMMp catalogue. \begin{theacknowledgments} This work has been supported by Funda\c c\~ao de Amparo \`a Pesquisa do Estado de S\~ao Paulo (FAPESP), Coordena\c c\~ao de Aperfei\c coamento de Pessoal de N\'ivel Superior (CAPES), Brazil and by Universit\'e Louis Pasteur, Strasbourg, France. \end{theacknowledgments} | 7 | 10 | 0710.5192 |
0710 | 0710.3816.txt | Spectral studies of quiescent emission and bursts of magnetar candidates using XMM-Newton, {\it Chandra} and {\it Swift} data are presented. Spectra of both the quiescent emission and the bursts for most magnetar candidates are reproduced by a photoelectrically absorbed two blackbody function (2BB). There is a strong correlation between lower and higher temperatures of 2BB ($kT_{\mathrm{LT}}$ and $kT_{\mathrm{HT}}$) for the magnetar candidates of which the spectra are well reproduced by 2BB. In addition, a square of radius for $kT_{\mathrm{LT}}$ ($R_{\mathrm{LT}}^2$) is well correlated with a square of radius for $kT_{\mathrm{HT}}$ ($R_{\mathrm{HT}}^2$). % A ratio $kT_{\mathrm LT}/kT_{\mathrm HT} \approx 0.4$ is nearly constant irrespective of objects and/or emission types (i.e., the quiescent emission and the bursts). This would imply a common emission mechanism among the magnetar candidates. The relation between the quiescent emission and the bursts might be analogous to a relation between microflares and solar flares of the sun. Three AXPs (4U\,0142$+$614, 1RXS\,J170849.0$-$400910 and 1E\,2259$+$586) seem to have an excess above $\sim$\,7\,keV which well agrees with a non-thermal hard component discovered by INTEGRAL. | Among peculiar celestial objects in the universe, a dense highly magnetized neutron star ($\rho \sim 10^{14}$\,g\,cm$^{-3}$ and $B \sim 10^{15}$\,G), so-called ``magnetar'' \citep{duncan1992, paczynski1992, thompson1995, thompson1996}, would be one of the most exotic objects. % Soft gamma repeaters (SGRs) and anomalous X-ray pulsars (AXPs) are well known as magnetar candidates. An apparent difference between the SGRs and the AXPs would be considered from their first detections. The SGRs were discovered as sporadically bursting objects, while the AXPs were regarded as peculiar pulsars with long spin periods. However, current observations unveil a lot of similarities between these objects. They have, for instance, long spin periods ($P \sim $ 5-12\,s) with spindown rates of $\dot{P} \sim 10^{-10}$-$10^{-13}$\,s\,s$^{-1}$, no signature of a companion star, a distribution around the galactic plane (two magnetar candidates are in other galaxies), quiescent soft X-ray emission. Several of these objects have non-thermal hard ($>20$\,keV) components, some are associated with supernova remnants (SNRs), and bursting activity is not confined to the SGRs but is observed in the AXPs as well. Considering these similarities, the SGRs and the AXPs should be classified into a common class of objects. So far, five SGRs (0501$+$4516, 0526$-$66, 1627$-$41, 1806$-$20 and 1900$+$14) are known \citep{woods2006, barthelmy2008} as well as three candidates, SGR\,1801$-$23 \citep{cline2000}, SGR\,1808$-$20 \citep{lamb2003} and SGR/GRB\,050925. SGR/GRB\,050925 was regarded as a gamma-ray burst (GRB) when first detected, but soon after was recognized as a new SGR \citep{holland2005}. On the other hand, ten AXPs (1E\,2259$+$586, 1E\,1048.1$-$5937, 4U\,0142$+$614, 1RXS\,J170849.0$-$400910, 1E\,1841$-$045, XTE\,J1810$-$197, AX\,J1845$-$0258, CXOU\,J010043.1$-$721134, CXOU\,J164710.2$-$455216 and 1E\,1547.0$-$5408) are known to date \citep{woods2006, dib2008} with one AXP candidate, AXP\,CXOU\,J160103.1$-$513353 \citep{park2006}. % A short burst from AXP\,CXOU\,J164710.2$-$455216 was detected by {\it Swift} BAT \citep{krimm2006} at 01:34:52 on 2006 September 21. The follow-up observations performed by {\it Swift} XRT found a remarkable result in which the quiescent emission of post-burst became 190 times brighter than that of pre-burst \citep{campana2006}. In addition to these objects, AX\,J1818.8$-$1559 discovered by ASCA \citep{sugizaki2001} recently exhibited a short burst \citep{mereghetti2007b} similar to those from the magnetar candidates. Therefore AX\,J1818.8$-$1559 could be a new SGR or AXP \citep{mereghetti2007b}. The most exciting phenomena among the magnetar candidates would be a sudden release of huge energy in rather short period, the so-called {\it giant flares} from the SGRs. They typically have a short intense spike which last less than 1\,s, and followed by a long pulsating tail which lasts a few hundred seconds. Their peak energy flux can be larger than $\sim10^{6}$ times Eddington luminosity. Theoretical studies suggested that the giant flares were triggered by a catastrophic deformation of the neutron star crust due to a torsion of the strong magnetic field (e.g., \cite{thompson2001}). Some different emission mechanisms have been proposed by several authors \citep{yamazaki2005, lyutikov2006, cea2006}. % In the past three decades, three giant flares were recorded. The first detection, from the source now known as SGR\,0526$-$66 in Large Magellanic Cloud (LMC), was made on March 5 in 1979 \citep{mazets1979, cline1980, evans1980, fenimore1996}. % The second one from SGR\,1900$+$14 was recorded on August 27 in 1998 \citep{hurley1999b, feroci1999, mazets1999, feroci2001, tanaka2007}. More recently, the most energetic giant flare from SGR\,1806$-$20 was observed on December 27 in 2004 \citep{cameron2005, gaensler2005, hurley2005, mazets2005, palmer2005, terasawa2005, tanaka2007}. The fluence of its initial intense spike with 600\,ms was evaluated to be $\sim 2$\,erg\,cm$^{-2}$ by the plasma particle detectors on the Geotail space probe \citep{terasawa2005}. Soft X-ray spectra of the quiescent emission of the SGRs and the AXPs were observed by a number of satellites. Although their spectral model is still under discussion, two two-component models are proposed. % One of them is a photoelectrically absorbed two blackbody function (2BB). Spectral parameters of 2BB are reported by some authors for the SGRs \citep{mereghetti2006a} and the AXPs \citep{tiengo2002, morii2003, gotthelf2004, gotthelf2005, halpern2005, tiengo2005, israel2006, gotthelf2007}. % Typical lower and higher temperatures are $\sim 0.5$\,keV and $\sim 1.4$\,keV, respectively. % The other model is a photoelectrically absorbed power law plus a blackbody (PL$+$BB). Some authors report spectral parameters of PL$+$BB for the SGRs \citep{marsden2001, kurkarni2003, mereghetti2005, mereghetti2006a, mereghetti2007a} and the AXPs \citep{morii2003, patel2003, rea2003, gotthelf2004, mereghetti2004, woods2004, tiengo2005, gavriil2006, israel2006}. % A typical power law index and a blackbody temperature are $\sim 3$ and $\sim 0.5$\,keV. % At present it is still unclear which model is more reliable or physically suitable. Recent observations by the International Gamma-Ray Astrophysics Laboratory (INTEGRAL) discovered a non-thermal hard component in the spectra of the quiescent emission above 20\,keV for 5 magnetar candidates \citep{molkov2005, gotz2006b, kuiper2006}. The non-thermal hard component is well reproduced by a power law model, $E^{-\Gamma}$, where $\Gamma$ ranges from 1.0-1.8, while the soft X-ray emission below $\sim12$\,keV, mentioned above, clearly indicates steeper power-law index of $\sim 3$ if the PL$+$BB is applied as the model spectrum. Hence, the non-thermal hard emission seen by INTEGRAL is a different component and presumably has a different origin than the soft X-ray emission. Since some magnetar candidates have two different emission mechanisms, there seems to be more complex physics than expected before. % Moreover, the non-thermal hard component shows pulsations for three AXPs, 1RXS\,J170849.0$-$400910, 4U\,0142$+$614 and 1E\,1841$-$045, through the INTEGRAL and RXTE observations \citep{kuiper2006}, which is related to a neutron star rotation, and hence there are particle acceleration processes in the vicinity of neutron stars \citep{kuiper2006}. If the energy source of the quiescent emission and the % short bursts is the magnetic field as thought to be, at least very similar physical process would govern both of them and their spectra could emerge alike. It is claimed based on High Energy Transient Explorer 2 (HETE-2) data that the most acceptable spectral model of the short bursts from two SGRs 1806¡Ý20 and 1900+14 is 2BB even though it should be regarded just as an empirical model (Nakagawa et al. 2007). It would be also preferred to represent spectra of quiescent emissions by 2BB rather than BB+PL for SGRs, and even for AXPs if it is the same class of object. In this paper, we present a comprehensive spectral study with 2BB for both the quiescent emission and the for the magnetar candidates. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% | The spectral studies using the photoelectrically absorbed two blackbody function (2BB) were presented for the quiescent emission and the burst of the magnetar candidates. % The spectra of the quiescent emission were well reproduced by a 2BB with some exceptions. The spectra of three AXPs (4U\,0142$+$614, 1RXS\,J170849.0$-$400910 and 1E\,2259$+$586) seem to have an excess which might be due to a non-thermal hard component discovered by INTEGRAL. % The spectrum of the burst from the SGR candidate SGR\,2013$+$34 was also well reproduced by 2BB. A strong linear correlations between $kT_{\mathrm{LT}}$ and $kT_{\mathrm{HT}}$ was found using 2BB spectra. The ratio $kT_{\mathrm{LT}}/kT_{\mathrm{HT}} \sim 0.4$ is almost constant irrespectively of the objects and/or emission types (burst or quiescent emission). The relationship between $R_{\mathrm{LT}}^2$ and $R_{\mathrm{HT}}^2$ seems to have a linear correlation. Considering these correlations, there seems to be a common emission mechanism among these objects, and between the quiescent emission and the burst. The relationship between the quiescent emission and the burst might be similar to the relationship between microflares and an ordinary solar flares of the sun. The quiescent emission might be due to very frequent small activity similar to the microflares. On the other hand, the burst might be due to a relatively large activity similar to the ordinary solar flare. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% % Acknowledgement \bigskip We would like to thank an anonymous referee for helpful comments and suggestions to improve our paper. This work is based on observations obtained with XMM-Newton, an ESA science mission with instruments and contributions directly funded by ESA Member States and NASA. % We would like to thank public data archive of {\it Chandra}. % This research has made use of software provided by the Chandra X-ray Center (CXC) in the application packages CIAO, ChIPS, and Sherpa. % We acknowledge the use of public data from the Swift data archive. % YEN is supported by the JSPS Research Fellowships for Young Scientists. This work is supported in part by a special postdoctoral researchers program in RIKEN. %\appendix %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% % Figures %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %\clearpage % \begin{figure} \begin{center} \FigureFile(80mm,129mm){figure1.eps} \end{center} \caption{Time history of the (a) flux in 2-10 keV in units of ergs cm$^{-2}$ s$^{-1}$, (b) photoelectric absorption $N_{\mathrm{H}}$ in units of cm$^{-2}$, (c) temperature of the lower blackbody $kT_{\mathrm{LT}}$ in units of keV, (d) radius of the lower blackbody $R_{\mathrm{LT}}$ in units of km, (e) temperature of the higher blackbody $kT_{\mathrm{HT}}$ in units of keV and (f) radius of the higher blackbody $R_{\mathrm{HT}}$ in units of km for the emissions of AXP\,CXOU\,J164710.2$-$455216 observed by XRT/{\it Swift}.}\label{fig:lc_persistent_cxou_j16} \end{figure} \begin{figure} \begin{center} \FigureFile(80mm,80mm){figure2.eps} \end{center} \caption{A schematic view of $\nu{F}_{\nu}$ spectra using 2BB$+$PL (a) and BB$+$2PL (b) for AXP\,4U\,0142$+$614. Spectral parameters of X-ray spectra (i.e., $\lesssim$12\,keV) are derived from our analyses using the XMM-Newton observation of 0112781101, while those of the non-thermal hard component (i.e., $\gtrsim$20\,keV) is obtained by INTEGRAL observations \citep{kuiper2006}. The circles denote data derived from our analyses using the XMM-Newton observation of 0112781101. The squares represent INTEGRAL observations taken from Fig.7 in \citet{kuiper2006} by eye. The dashed, dot-dash, dotted lines in (a) show the $kT_{\mathrm{LT}}$ component, the $kT_{\mathrm{HT}}$ component and the PL component for the hard spectrum, respectively. Those lines in (b) show the $kT$ component, the PL component for an X-ray spectrum and the PL component for the hard spectrum, respectively.}\label{fig:eeufspec_summary} \end{figure} \begin{figure} \begin{center} \FigureFile(80mm,83mm){figure3.eps} \end{center} \caption{Relationship between the 2BB temperatures $kT_{\mathrm{LT}}$ and $kT_{\mathrm{HT}}$. The triangles and squares denote the previous work on the bursts \citep{feroci2004, olive2004, gotz2006a, nakagawa2007} and the quiescent emission \citep{morii2003, gotthelf2004, gotthelf2005, tiengo2005, mereghetti2006a}, respectively. The circles and stars denote our work on the bursts and the quiescent emission, respectively. The line represents the best-fit power law model.}\label{fig:bb_relation} \end{figure} \begin{figure} \begin{center} \FigureFile(80mm,80mm){figure4.eps} \end{center} \caption{Relationship between the square of the blackbody radii $R^2_{\mathrm{LT}}$ and $R^2_{\mathrm{HT}}$. The triangles and squares denote the previous work on the the bursts \citep{olive2004, nakagawa2007} and the quiescent emission \citep{morii2003, tiengo2005, mereghetti2006a}, respectively. The stars denote our work on the quiescent emission. The solid line shows a ratio of $R_{\mathrm{HT}}^{2}$ to $R_{\mathrm{LT}}^{2}$ of 0.01.}\label{fig:radius_relation} \end{figure} \clearpage \begin{figure} \begin{center} \FigureFile(160mm,80mm){figure5.eps} \end{center} \caption{{\it Left:} Relationship between the lower temperature of 2BB $kT_{\mathrm{LT}}$ and the square of the blackbody radii of 2BB $R_{\mathrm{LT}}^2$. {\it Right:} Relationship between the higher temperature of 2BB $kT_{\mathrm{HT}}$ and the square of the blackbody radii $R_{\mathrm{HT}}^2$. The dotted and dashed lines correspond to bolometric fluences of $10^{-8}$ and $10^{-9}$\,ergs\,cm$^{-2}$, respectively.}\label{fig:kt_rsquared_relation} \end{figure} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% % Tables %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \clearpage \begin{table*} \small \caption{A summary of magnetar candidates which were employed in our study.}\label{tab:magnetar_list} \begin{center} \begin{tabular}{lllll} \hline\hline Object\footnotemark[$*$] & Satellite/Instrument\footnotemark[$\dagger$] & Period\footnotemark[$\ddagger$] & Distance\footnotemark[$\S$] & Ref.\footnotemark[$\|$] \\ \hline SGR\,0526$-$66 & {\it Chandra} ACIS & 2000-2001 & 50 & (1) \\ SGR\,1627$-$41 & XMM-Newton EPIC, {\it Chandra} ACIS & 2001-2005 & 11 & (2), (3), (4) \\ SGR\,1806$-$20 & XMM-Newton EPIC, {\it Chandra} ACIS & 2000-2005 & 15\footnotemark[$\#$] & (5), (6) \\ SGR\,2013$+$34 & {\it Swift} BAT & 2005 & 10 & (7), (8) \\ SGR\,1819$-$16 & XMM-Newton EPIC & 2003 & 10 & \\ AXP\,CXOU\,J010043.1$-$721134 & XMM-Newton EPIC, {\it Chandra} ACIS & 2000-2005 & 57 & (9), (10) \\ AXP\,4U\,0142$+$614 & XMM-Newton EPIC & 2002-2004 & 3 & (11) \\ AXP\,CXOU\,J164710.2$-$455216 & {\it Swift} BAT, {\it Swift} XRT, {\it Chandra} ACIS & 2005-2007 & 5 & (12), (13), (14), (15) \\ AXP\,1RXS\,J170849.0$-$400910 & XMM-Newton EPIC & 2003 & 5 & (16), (17) \\ AXP\,1E\,1841$-$045 & XMM-Newton EPIC & 2002 & 7\footnotemark[$**$] & \\ AXP\,1E\,2259$+$586 & XMM-Newton EPIC & 2002-2005 & 3 & (18) \\ \hline \multicolumn{5}{@{}l@{}}{\hbox to 0pt{\parbox{180mm}{\footnotesize \footnotemark[$*$] Object name of magnetar candidates (SGR\,2013$+$34 denotes SGR candidate SGR/GRB\,050925). \par\noindent \footnotemark[$\dagger$] Instrument and satellite names from which obtained the data used in our analysis. \par\noindent \footnotemark[$\ddagger$] Interval which these observations were performed. \par\noindent \footnotemark[$\S$] Distances to each object in units of kpc (see \cite{woods2006} and references there in). \par\noindent \footnotemark[$\|$] (1) \citet{kurkarni2003}; (2) \citet{kouveliotou2003}; (3) \citet{wachter2004}; (4) \citet{mereghetti2006b}; (5) \citet{kaplan2002}; (6) \citet{mereghetti2005}; (7) \citet{holland2005}; (8) \citet{markwardt2005}; (9) \citet{lamb2002}; (10) \citet{majid2004}; (11) \citet{guver2007}; (12) \citet{krimm2006}; (13) \citet{campana2006}; (14) \citet{muno2006}; (15) \citet{muno2007}; (16) \citet{oosterbroek2004}; (17) \citet{rea2007}; (18) \citet{woods2004} \par\noindent \footnotemark[$\#$] The latest distance estimate is $d = 8.7$\,kpc \citep{bibby2008}. \par\noindent \footnotemark[$**$] The latest distance estimate is $d = 8.5$\,kpc \citep{tian2008}. }\hss}} \end{tabular} \end{center} \end{table*} \begin{table*} \scriptsize \caption{XMM-Newton observations of the quiescent emissions of the SGRs.}\label{tab:sgr_obs_xmm} \begin{center} \begin{tabular}{lllllllllllll} \hline\hline Object\footnotemark[$*$] & ObsID\footnotemark[$\dagger$] & \multicolumn{2}{l}{Observation Date (MJD)\footnotemark[$\ddagger$]} & \multicolumn{3}{l}{Observation Mode\footnotemark[$\S$]} & \multicolumn{3}{l}{Exposure Time (ks)\footnotemark[$\|$]} & \multicolumn{3}{l}{Source/Background Radii} \\ & & Start & End & pn & MOS1 & MOS2 & pn & MOS1 & MOS2 & pn & MOS1 & MOS2 \\ \hline 1627$-$41 & 0204500201 & 53051.590 & 53051.992 & Full & Full & Full & 15.89 & 20.03 & 20.17 & $\timeform{10''}$/$\timeform{10''}$ & $\timeform{10''}$/$\timeform{10''}$ & $\timeform{10''}$/$\timeform{10''}$ \\ 1627$-$41 & 0204500301 & 53252.750 & 53253.131 & Full & Full & Full & 27.06 & 32.00 & 32.10 & $\timeform{10''}$/$\timeform{10''}$ & $\timeform{10''}$/$\timeform{10''}$ & $\timeform{10''}$/$\timeform{10''}$ \\ 1627$-$41 & 0202560101 & 53270.677 & 53271.281 & Small & P-W2 & P-W2 & 26.66 & 36.19 & 49.38 & $\timeform{10''}$/$\timeform{10''}$ & $\timeform{10''}$/$\timeform{10''}$ & $\timeform{10''}$/$\timeform{10''}$ \\ 1806$-$20 & 0148210101 & 52732.566 & 52733.209 & Full & P-W3 & P-W3 & 4.84 & 5.65 & 5.62 & $\timeform{32''}$/$\timeform{32''}$ & $\timeform{28''}$/$\timeform{28''}$ & $\timeform{28''}$/$\timeform{28''}$ \\ 1806$-$20 & 0148210401 & 52919.404 & 52919.663 & Full & P-W3 & P-W3 & 7.61 & 7.07 & 7.23 & $\timeform{32''}$/$\timeform{32''}$ & $\timeform{28''}$/$\timeform{28''}$ & $\timeform{28''}$/$\timeform{28''}$ \\ 1806$-$20 & 0205350101 & 53254.377 & 53254.978 & Small & P-W3 & P-W3 & 30.21 & 39.14 & 39.56 & $\timeform{32''}$/$\timeform{32''}$ & $\timeform{28''}$/$\timeform{28''}$ & $\timeform{28''}$/$\timeform{28''}$ \\ 1806$-$20 & 0164561101 & 53284.706 & 53284.925 & Small & Fast-U & Fast-U & 11.54 & $t$ & $t$ & $\timeform{32''}$/$\timeform{32''}$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ \\ 1806$-$20 & 0164561301 & 53436.348 & 53436.636 & Small & Fast-U & Full & 7.37 & $t$ & 5.68 & $\timeform{32''}$/$\timeform{32''}$ & $\cdot\cdot\cdot$ & $\timeform{28''}$/$\timeform{28''}$ \\ 1806$-$20 & 0164561401 & 53647.427 & 53647.809 & Small & Fast-U & Full & 22.11 & $t$ & 28.68 & $\timeform{32''}$/$\timeform{32''}$ & $\cdot\cdot\cdot$ & $\timeform{28''}$/$\timeform{28''}$ \\ 2013$+$34 & 0212481201 & 53655.026 & 53655.334 & Full & Full & Full & 22.18 & 25.27 & 25.27 & $\timeform{32''}$/$\timeform{32''}$ & $\timeform{28''}$/$\timeform{28''}$ & $\timeform{28''}$/$\timeform{28''}$ \\ 1819$-$16 & 0152834501 & 52726.191 & 52726.310 & Full & Full & Full & 3.3 & 4.3 & 4.8 & $\timeform{32''}$/$\timeform{80''}$\footnotemark[$\#$] & $\timeform{28''}$/$\timeform{70''}$\footnotemark[$\#$] & $\timeform{28''}$/$\timeform{70''}$\footnotemark[$\#$] \\ \hline \multicolumn{13}{@{}l@{}}{\hbox to 0pt{\parbox{180mm}{\footnotesize \footnotemark[$*$] Object name of the SGRs (2013$+$34 denotes SGR candidate SGR/GRB\,050925 and 1819$-$16 denotes SGR candidate AX\,J1818.8$-$1559). \par\noindent \footnotemark[$\dagger$] XMM-Newton observation ID. \par\noindent \footnotemark[$\ddagger$] Start and end time of observations. \par\noindent \footnotemark[$\S$] Observation mode for each instrument; full-window mode (Full), small-window mode (Small), partial-w2 mode (P-W2), partial-w3 mode (P-W3) and fast-uncompressed mode (Fast-U). \par\noindent \footnotemark[$\|$] Net exposure time for each instrument. $t$ denotes the data sets obtained by the MOS cameras in timing mode and not utilized. \par\noindent \footnotemark[$\#$] The background regions were extracted from an annular region whose center was the source position. The first values are source radii, and the inner radii of the background regions. The second values are outer radii of the background regions. }\hss}} \end{tabular} \end{center} \end{table*} \begin{table*} \small \caption{{\it Chandra} observations of the quiescent emissions of the SGRs.}\label{tab:sgr_obs_cxo} \begin{center} \begin{tabular}{lllllll} \hline\hline Object\footnotemark[$*$] & ObsID\footnotemark[$\dagger$] & \multicolumn{2}{l}{Observation Date (MJD)\footnotemark[$\ddagger$]} & Observation Mode\footnotemark[$\S$] & Exposure Time (ks)\footnotemark[$\|$] & Source/Background Radii \\ & & Start & End & & & \\ \hline 0526$-$66 & 747 & 51547.017 & 51547.539 & FAINT & 39.86 & $\timeform{1''}$/$\timeform{1''}$ \\ 0526$-$66 & 1957 & 52152.937 & 52153.566 & FAINT & 48.45 & $\timeform{1''}$/$\timeform{1''}$ \\ 1627$-$41 & 1981 & 52182.205 & 52182.803 & FAINT & 48.93 & $\timeform{2''}$/$\timeform{2''}$ \\ 1627$-$41 & 3877 & 52722.169 & 52722.494 & VFAINT & 25.67 & $\timeform{2''}$/$\timeform{2''}$ \\ \hline \multicolumn{7}{@{}l@{}}{\hbox to 0pt{\parbox{180mm}{\footnotesize \footnotemark[$*$] SGR names. \par\noindent \footnotemark[$\dagger$] {\it Chandra} observation ID. \par\noindent \footnotemark[$\ddagger$] Start and end time of the observations. \par\noindent \footnotemark[$\S$] FAINT and VFAINT denote the imaging mode, and CC33\_FAINT denotes the timing mode. \par\noindent \footnotemark[$\|$] Net exposure time. }\hss}} \end{tabular} \end{center} \end{table*} \begin{table*} \scriptsize \caption{XMM-Newton observations of the quiescent emissions of the AXPs.}\label{tab:axp_obs_xmm} \begin{center} \begin{tabular}{lllllllllllll} \hline\hline Object\footnotemark[$*$] & ObsID\footnotemark[$\dagger$] & \multicolumn{2}{l}{Observation Date (MJD)\footnotemark[$\ddagger$]} & \multicolumn{3}{l}{Observation Mode\footnotemark[$\S$]} & \multicolumn{3}{l}{Exposure Time (ks)\footnotemark[$\|$]} & \multicolumn{3}{l}{Source/Background Radii} \\ & & Start & End & pn & MOS1 & MOS2 & pn & MOS1 & MOS2 & pn & MOS1 & MOS2 \\ \hline 0100$-$721 & 0110000201 & 51834.626 & 51834.867 & E-Full & Full & Full & 20.81 & 14.62 & $g$ & $\timeform{32''}$/$\timeform{32''}$ & $\timeform{28''}$/$\timeform{28''}$ & $\timeform{28''}$/$\timeform{28''}$ \\ 0100$-$721 & 0018540101 & 52233.983 & 52234.303 & Full & Full & Full & 21.16 & 25.73 & $g$ & $\timeform{32''}$/$\timeform{32''}$ & $\timeform{28''}$/$\timeform{28''}$ & $\timeform{28''}$/$\timeform{28''}$ \\ 0142$+$614 & 0112781101 & 52663.920 & 52663.995 & Small & Fast-U & Fast-U & 4.18 & $t$ & $t$ & $\timeform{32''}$/$\timeform{32''}$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ \\ 1708$-$400 & 0148690101 & 52879.906 & 52880.426 & Full & P-W3 & P-W3 & 26.88 & $p$ & $p$ & $\timeform{20''}$/$\timeform{20''}$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ \\ 1841$-$045 & 0013340101 & 52552.122 & 52552.192 & Large & Full & Full & 2.34 & 3.54 & 3.57 & $\timeform{12''}$/$\timeform{12''}$ & $\timeform{12''}$/$\timeform{12''}$ & $\timeform{12''}$/$\timeform{12''}$ \\ 1841$-$045 & 0013340201 & 52554.115 & 52554.193 & Large & Full & Full & 4.37 & 6.30 & 6.30 & $\timeform{12''}$/$\timeform{12''}$ & $\timeform{12''}$/$\timeform{12''}$ & $\timeform{12''}$/$\timeform{12''}$ \\ 2259$+$586 & 0038140101 & 52436.378 & 52436.986 & Small & Full & Full & 30.63 & $p$ & $p$ & $\timeform{32''}$/$\timeform{32''}$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ \\ 2259$+$586 & 0155350301 & 52446.400 & 52446.759 & Small & P-W2 & Full & 17.65 & $p$ & $p$ & $\timeform{32''}$/$\timeform{32''}$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ \\ 2259$+$586 & 0203550701 & 53579.965 & 53580.030 & Small & P-W2 & Fast-U & 2.66 & $p$ & $t$ & $\timeform{32''}$/$\timeform{32''}$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ \\ \hline \multicolumn{13}{@{}l@{}}{\hbox to 0pt{\parbox{180mm}{\footnotesize \footnotemark[$*$] Object name of the AXPs; CXOU\,J010043.1$-$721134 (0100$-$721), 4U\,0142$+$614 (0142$+$614), 1RXS\,J170849.0$-$400910 (1708$-$400), 1E\,1841$-$045 and 1E\,2259$+$586. \par\noindent \footnotemark[$\dagger$] XMM-Newton observation ID. \par\noindent \footnotemark[$\ddagger$] Start and end time of the observations. \par\noindent \footnotemark[$\S$] Observation mode for each instrument; extended full-window mode (E-Full), Full-window mode (Full), small-window mode (Small), fast-uncompressed mode (Fast-U), fast-timing mode (Fast-T), partial-w3 mode (P-W3), large-window mode (Large) and partial-w2 mode (P-W2). \par\noindent \footnotemark[$\|$] Net exposure time for each instrument. $g$ denotes that the source fell on a gap of the CCD chips, $t$ denotes observations in timing mode, and $p$ denotes that the data sets are affected by a photon pile-up. These data sets were not utilized. }\hss}} \end{tabular} \end{center} \end{table*} \begin{table*} \small \caption{{\it Chandra} observations of the quiescent emissions of the AXPs.}\label{tab:axp_obs_cxo} \begin{center} \begin{tabular}{lllllll} \hline\hline Object\footnotemark[$*$] & ObsID\footnotemark[$\dagger$] & \multicolumn{2}{l}{Observation Date (MJD)\footnotemark[$\ddagger$]} & Observation Mode\footnotemark[$\S$] & Exposure Time (ks)\footnotemark[$\|$] & Source/Background Radii \\ & & Start & End & & & \\ \hline 0100$-$721 & 1881 & 52044.080 & 52045.261 & FAINT & 98.67 & $\timeform{11''}$/$\timeform{10''}$\footnotemark[$\#$] \\ 0100$-$721 & 4616 & 53031.791 & 53032.009 & VFAINT & 15.56 & $\timeform{2''}$/$\timeform{2''}$ \\ 0100$-$721 & 4617 & 53032.189 & 53032.399 & VFAINT & 15.27 & $\timeform{2''}$/$\timeform{2''}$ \\ 0100$-$721 & 4618 & 53033.904 & 53034.130 & VFAINT & 15.00 & $\timeform{2''}$/$\timeform{2''}$ \\ 0100$-$721 & 4619 & 53042.806 & 53043.023 & VFAINT & 15.04 & $\timeform{2''}$/$\timeform{2''}$ \\ 0100$-$721 & 4620 & 53089.185 & 53089.395 & VFAINT & 15.22 & $\timeform{2''}$/$\timeform{2''}$ \\ 1647$-$455 & 6283 & 53512.860 & 53513.102 & FAINT & 18.81 & $\timeform{2''}$/$\timeform{2''}$ \\ 1647$-$455 & 5411 & 53539.673 & 53540.141 & FAINT & 38.47 & $\timeform{2''}$/$\timeform{2''}$ \\ \hline \multicolumn{7}{@{}l@{}}{\hbox to 0pt{\parbox{180mm}{\footnotesize \footnotemark[$*$] Object name of the AXPs; CXOU\,J010043.1$-$721134 (0100$-$721) and CXOU\,J164710.2$-$455216 (1647$-$455). \par\noindent \footnotemark[$\dagger$] {\it Chandra} observation ID. \par\noindent \footnotemark[$\ddagger$] Start and end time of observations. \par\noindent \footnotemark[$\S$] FAINT and VFAINT denote the imaging mode. \par\noindent \footnotemark[$\|$] Net exposure time. \par\noindent \footnotemark[$\#$] Since the source fell on an off-axis CCD chip, the source region was extracted from an elliptical region with major and minor axes of $\timeform{11''}$ and $\timeform{10''}$, respectively. }\hss}} \end{tabular} \end{center} \end{table*} \begin{table*} \small \caption{{\it Swift} observations of the quiescent emission and the bursts of AXP\,CXOU\,J164710.2$-$455216.}\label{tab:axp_j1647_obs_swift} \begin{center} \begin{tabular}{lllllllll} \hline\hline SeqNum\footnotemark[$*$] & \multicolumn{4}{l}{Observation Date (MJD)\footnotemark[$\dagger$]} & \multicolumn{2}{l}{Exposure Time (ks)\footnotemark[$\ddagger$]} & \multicolumn{2}{l}{Source/Background Radii} \\ & Start (WT) & End (WT) & Start (PC) & End (PC) & WT & PC & WT & PC \\ \hline 00230341000\footnotemark[$\S$] & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ \\ 00030806001 & 53999.604 & 53999.924 & 53999.604 & 53999.936 & 1.92 & 7.74 & $\timeform{36''}\times\timeform{18''}$/$\timeform{36''}\times\timeform{18''}$\footnotemark[$\|$] & $\timeform{30''}$/$\timeform{30''}$ \\ 00030806002 & 54000.610 & 54000.619 & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & 0.77 & $\cdot\cdot\cdot$ & $\timeform{36''}\times\timeform{18''}$/$\timeform{36''}\times\timeform{18''}$\footnotemark[$\|$] & $\cdot\cdot\cdot$ \\ 00030806003 & 54000.819 & 54001.212 & 54000.819 & 54001.073 & 4.91 & 1.84 & $\timeform{36''}\times\timeform{18''}$/$\timeform{36''}\times\timeform{18''}$\footnotemark[$\|$] & $\timeform{30''}$/$\timeform{30''}$ \\ 00030806004 & 54004.276 & 54004.483 & 54004.343 & 54004.489 & 1.25 & 2.48 & $\timeform{36''}\times\timeform{18''}$/$\timeform{36''}\times\timeform{18''}$\footnotemark[$\|$] & $\timeform{30''}$/$\timeform{30''}$ \\ 00030806006 & 54010.461 & 54010.716 & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & 1.98 & $\cdot\cdot\cdot$ & $\timeform{36''}\times\timeform{18''}$/$\timeform{36''}\times\timeform{18''}$\footnotemark[$\|$] & $\cdot\cdot\cdot$ \\ 00030806007 & 54011.515 & 54011.594 & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & 2.03 & $\cdot\cdot\cdot$ & $\timeform{36''}\times\timeform{18''}$/$\timeform{36''}\times\timeform{18''}$\footnotemark[$\|$] & $\cdot\cdot\cdot$ \\ 00030806008 & 54014.001 & 54014.077 & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & 2.16 & $\cdot\cdot\cdot$ & $\timeform{36''}\times\timeform{18''}$/$\timeform{36''}\times\timeform{18''}$\footnotemark[$\|$] & $\cdot\cdot\cdot$ \\ 00030806009 & 54017.746 & 54017.954 & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & 3.52 & $\cdot\cdot\cdot$ & $\timeform{36''}\times\timeform{18''}$/$\timeform{36''}\times\timeform{18''}$\footnotemark[$\|$] & $\cdot\cdot\cdot$ \\ 00030806010 & 54018.009 & 54018.094 & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & 2.83 & $\cdot\cdot\cdot$ & $\timeform{36''}\times\timeform{18''}$/$\timeform{36''}\times\timeform{18''}$\footnotemark[$\|$] & $\cdot\cdot\cdot$ \\ 00030806011 & 54023.237 & 54023.517 & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & 5.62 & $\cdot\cdot\cdot$ & $\timeform{36''}\times\timeform{18''}$/$\timeform{36''}\times\timeform{18''}$\footnotemark[$\|$] & $\cdot\cdot\cdot$ \\ 00030806012 & 54029.125 & 54029.342 & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & 5.52 & $\cdot\cdot\cdot$ & $\timeform{36''}\times\timeform{18''}$/$\timeform{36''}\times\timeform{18''}$\footnotemark[$\|$] & $\cdot\cdot\cdot$ \\ 00030806013 & 54035.677 & 54035.825 & 54035.678 & 54035.774 & 2.82 & 0.42 & $\timeform{36''}\times\timeform{18''}$/$\timeform{36''}\times\timeform{18''}$\footnotemark[$\|$] & $\timeform{30''}$/$\timeform{30''}$ \\ 00030806014 & 54119.174 & 54119.253 & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & 2.06 & $\cdot\cdot\cdot$ & $\timeform{36''}\times\timeform{18''}$/$\timeform{36''}\times\timeform{18''}$\footnotemark[$\|$] & $\cdot\cdot\cdot$ \\ 00030806015 & 54122.050 & 54122.267 & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & 3.82 & $\cdot\cdot\cdot$ & $\timeform{36''}\times\timeform{18''}$/$\timeform{36''}\times\timeform{18''}$\footnotemark[$\|$] & $\cdot\cdot\cdot$ \\ \hline \multicolumn{9}{@{}l@{}}{\hbox to 0pt{\parbox{180mm}{\footnotesize \footnotemark[$*$] {\it Swift} sequence number. \par\noindent \footnotemark[$\dagger$] Start and end time of the observations for each mode (WT denotes window timing mode, and PC denotes photon counting mode). \par\noindent \footnotemark[$\ddagger$] Net exposure time for each observation mode. \par\noindent \footnotemark[$\S$] The observation of a burst. \par\noindent \footnotemark[$\|$] The source and background regions were extracted from a rectangle region. Two background regions are utilized near both sides of the source region. }\hss}} \end{tabular} \end{center} \end{table*} \begin{table*} \caption{Spectral parameters of the quiescent emissions of the SGRs observed by XMM-Newton.}\label{tab:sgr_spc_xmm} \begin{center} \begin{tabular}{lllllllll} \hline\hline Object\footnotemark[$*$] & ObsID\footnotemark[$\dagger$] & $N_{\mathrm{H}}$\footnotemark[$\ddagger$] & $kT_{\mathrm{LT}}$\footnotemark[$\S$] & $R_{\mathrm{LT}}$\footnotemark[$\|$] & $kT_{\mathrm{HT}}$\footnotemark[$\S$] & $R_{\mathrm{HT}}$\footnotemark[$\|$] & $F$\footnotemark[$\#$] & $\chi^{2}$ (d.o.f.) \\ & & (10$^{22}$ cm$^{-2}$) & (keV) & (km) & (keV) & (km) & & \\ \hline 1627$-$41 & 0204500201 & 15.98$_{-7.60}^{+15.51}$ & 0.58$_{-0.25}^{+0.35}$ & $<$19.25 & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & $\sim$\,0.05 & 44 (42) \\ 1627$-$41 & 0204500301 & 7.53$_{-3.56}^{+6.19}$ & 0.85$_{-0.22}^{+0.29}$ & $<$0.62 & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & 0.08$\pm0.06$ & 29 (65) \\ 1627$-$41 & 0202560101 & 9.00$_{-3.89}^{+6.71}$ & 0.94$_{-0.24}^{+0.31}$ & $<$0.42 & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & 0.06$\pm0.04$ & 54 (47) \\ 1806$-$20 & 0148210101 & 5.18$_{-0.73}^{+0.92}$ & 0.84$_{-0.17}^{+0.23}$ & 1.64$_{-0.53}^{+1.04}$ & 2.62$_{-0.38}^{+0.97}$ & 0.28$_{-0.12}^{+0.10}$ & 11.05$\pm1.92$ & 312 (295) \\ 1806$-$20 & 0148210401 & 5.87$_{-0.55}^{+0.63}$ & 0.85$_{-0.11}^{+0.12}$ & 1.97$_{-0.43}^{+0.65}$ & 3.39$_{-0.58}^{+1.2}$ & 0.20$\pm0.07$ & 12.29$\pm2.4$ & 506 (459) \\ 1806$-$20 & 0205350101 & 5.75$_{-0.19}^{+0.20}$ & 0.96$\pm0.05$ & 2.06$_{-0.17}^{+0.20}$ & 3.19$_{-0.21}^{+0.28}$ & 0.32$\pm0.04$ & 25.43$\pm0.52$ & 2243 (2224) \\ 1806$-$20 & 0164561101 & 5.43$_{-0.35}^{+0.39}$ & 1.02$\pm0.1$ & 1.94$_{-0.27}^{+0.37}$ & 3.92$_{-0.69}^{+1.46}$ & 0.23$\pm0.08$ & 24.65$\pm2.74$ & 737 (741) \\ 1806$-$20 & 0164561301 & 5.64$_{-0.43}^{+0.50}$ & 0.96$\pm0.1$ & 1.98$_{-0.33}^{+0.46}$ & 3.67$_{-0.68}^{+1.54}$ & 0.22$\pm0.08$ & 18.91$\pm4.21$ & 654 (531) \\ 1806$-$20 & 0164561401 & 5.91$_{-0.31}^{+0.34}$ & 0.87$\pm0.06$ & 2.03$_{-0.25}^{+0.31}$ & 3.27$_{-0.34}^{+0.48}$ & 0.22$\pm0.04$ & 13.30$\pm0.5$ & 1026 (1069) \\ 2013$+$34 & 0212481201 & 0.29$_{-0.13}^{+0.15}$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & 0.13$_{-0.02}^{+0.02}$ & $<$7.54 & 20.78$\pm20.12$ & 57 (64) \\ 1819$-$16 & 0152834501 & $1.6_{-0.5}^{+0.7}$ & $1.9_{-0.2}^{+0.3}$ & $0.11_{-0.02}^{+0.03}$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & 1.3$\pm0.1$ & 59 (74) \\ \hline \multicolumn{9}{@{}l@{}}{\hbox to 0pt{\parbox{180mm}{\footnotesize \footnotemark[$*$] Object name of the SGRs (2013$+$34 denotes a SGR candidate SGR/GRB\,050925). \par\noindent \footnotemark[$\dagger$] XMM-Newton observation ID. \par\noindent \footnotemark[$\ddagger$] $N_{\mathrm{H}}$ denotes the column density with 90 \% confidence level errors. \par\noindent \footnotemark[$\S$] $kT_{\mathrm{LT}}$ and $kT_{\mathrm{HT}}$ denote the blackbody temperatures with 90 \% confidence level errors. \par\noindent \footnotemark[$\|$] $R_{\mathrm{LT}}$ and $R_{\mathrm{HT}}$ denote the emission radii with 90 \% confidence level errors. \par\noindent \footnotemark[$\#$] $F$ denotes a flux in the energy range 2-10 keV in units of $10^{-12}$ ergs cm$^{-2}$ s$^{-1}$ with 68 \% confidence level errors. }\hss}} \end{tabular} \end{center} \end{table*} \begin{table*} \caption{Spectral parameters of the quiescent emissions of the SGRs observed by {\it Chandra}.}\label{tab:sgr_spc_cxo} \begin{center} \begin{tabular}{lllllllll} \hline\hline Object\footnotemark[$*$] & ObsID\footnotemark[$\dagger$] & $N_{\mathrm{H}}$\footnotemark[$\ddagger$] & $kT_{\mathrm{LT}}$\footnotemark[$\S$] & $R_{\mathrm{LT}}$\footnotemark[$\|$] & $kT_{\mathrm{HT}}$\footnotemark[$\S$] & $R_{\mathrm{HT}}$\footnotemark[$\|$] & $F$\footnotemark[$\#$] & $\chi^{2}$ (d.o.f.) \\ & & (10$^{22}$ cm$^{-2}$) & (keV) & (km) & (keV) & (km) & & \\ \hline 0526$-$66 & 747 & 0.20$_{-0.02}^{+0.03}$ & 0.36$\pm0.03$ & 10.87$_{-1.26}^{+1.66}$ & 0.98$_{-0.14}^{+0.22}$ & 1.06$_{-0.37}^{+0.46}$ & 0.49$\pm0.06$ & 180 (181) \\ 0526$-$66 & 1957 & 0.26$_{-0.04}^{+0.05}$ & 0.30$\pm0.04$ & 15.08$_{-2.87}^{+4.93}$ & 0.70$_{-0.07}^{+0.11}$ & 2.24$_{-0.72}^{+0.85}$ & 0.39$\pm0.06$ & 176 (185) \\ 1627$-$41 & 1981 & 8.47$_{-4.86}^{+6.41}$ & 0.89$_{-0.28}^{+0.57}$ & $<$0.67 & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & $<$0.07 & 36 (34) \\ 1627$-$41 & 3877 & 17.34$_{-7.36}^{+9.89}$ & 0.42$_{-0.13}^{+0.2}$ & 3.56$_{-2.99}^{+33.39}$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & $<$0.06 & 22 (21) \\ \hline \multicolumn{9}{@{}l@{}}{\hbox to 0pt{\parbox{180mm}{\footnotesize \footnotemark[$*$] Object name of the SGRs. \par\noindent \footnotemark[$\dagger$] {\it Chandra} observation ID. \par\noindent \footnotemark[$\ddagger$] $N_{\mathrm{H}}$ denotes the column density with 90 \% confidence level errors. \par\noindent \footnotemark[$\S$] $kT_{\mathrm{LT}}$ and $kT_{\mathrm{HT}}$ denote the blackbody temperatures with 90 \% confidence level errors. \par\noindent \footnotemark[$\|$] $R_{\mathrm{LT}}$ and $R_{\mathrm{HT}}$ denote the emission radii with 90 \% confidence level errors. \par\noindent \footnotemark[$\#$] $F$ denotes the flux in the 2-10\,keV band in units of $10^{-12}$ ergs cm$^{-2}$ s$^{-1}$ with 68 \% confidence level errors. }\hss}} \end{tabular} \end{center} \end{table*} \begin{table*} \caption{Spectral parameters of a burst of SGR\,2013$+$34 (GRB/SGR\,050925) observed by {\it Swift}.}\label{tab:sgr_050925_spc_swift} \begin{center} \begin{tabular}{llllllll} \hline\hline SeqNum\footnotemark[$*$] & $N_{\mathrm{H}}$\footnotemark[$\dagger$] & $kT_{\mathrm{LT}}$\footnotemark[$\ddagger$] & $R_{\mathrm{LT}}$\footnotemark[$\S$] & $kT_{\mathrm{HT}}$\footnotemark[$\ddagger$] & $R_{\mathrm{HT}}$\footnotemark[$\S$] & $F$\footnotemark[$\|$] & $\chi^{2}$ (d.o.f.) \\ & (10$^{22}$ cm$^{-2}$) & (keV) & (km) & (keV) & (km) & & \\ \hline 00156838000 & $\cdot\cdot\cdot$ & 6.6$_{-3.9}^{+5.6}$ & 3.1$_{-1.7}^{+3.7}$ & 20$_{-4}^{+15}$ & $\gtrsim$0.2 & 0.81$\pm0.28$ & 27 (25) \\ \hline \multicolumn{8}{@{}l@{}}{\hbox to 0pt{\parbox{180mm}{\footnotesize \footnotemark[$*$] {\it Swift} sequence number. \par\noindent \footnotemark[$\dagger$] $N_{\mathrm{H}}$ denotes the column density with 90 \% confidence level errors. \par\noindent \footnotemark[$\ddagger$] $kT_{\mathrm{LT}}$ and $kT_{\mathrm{HT}}$ denote the blackbody temperatures with 90 \% confidence level errors. \par\noindent \footnotemark[$\S$] $R_{\mathrm{LT}}$ and $R_{\mathrm{HT}}$ denote the emission radii with 90 \% confidence level errors. \par\noindent \footnotemark[$\|$] $F$ denotes a flux in the energy range 15-150 keV in units of $10^{-6}$ ergs cm$^{-2}$ s$^{-1}$ with 68 \% confidence level errors. }\hss}} \end{tabular} \end{center} \end{table*} \begin{table*} \small \caption{Spectral parameters of the quiescent emissions of the AXPs observed by XMM-Newton.}\label{tab:axp_spc_xmm} \begin{center} \begin{tabular}{lllllllll} \hline\hline Object\footnotemark[$*$] & ObsID\footnotemark[$\dagger$] & $N_{\mathrm{H}}$\footnotemark[$\ddagger$] & $kT_{\mathrm{LT}}$\footnotemark[$\S$] & $R_{\mathrm{LT}}$\footnotemark[$\|$] & $kT_{\mathrm{HT}}$\footnotemark[$\S$] & $R_{\mathrm{HT}}$\footnotemark[$\|$] & $F$\footnotemark[$\#$] & $\chi^{2}$ (d.o.f.) \\ & & (10$^{22}$ cm$^{-2}$) & (keV) & (km) & (keV) & (km) & & \\ \hline 0100$-$721 & 0110000201 & $\lesssim0.06$ & 0.39$\pm0.02$ & 6.94$_{-0.67}^{+1.09}$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & 0.09$\pm0.01$ & 264 (275) \\ 0100$-$721 & 0018540101 & $\lesssim0.15$ & 0.29$\pm0.06$ & 11.64$_{-3.35}^{+7.47}$ & 0.64$_{-0.11}^{+0.25}$ & 1.86$_{-1.08}^{+1.17}$ & 0.13$\pm0.09$ & 97 (129) \\ 0142$+$614 & 0112781101 & 0.53$\pm0.01$ & 0.36$\pm0.01$ & 9.38$_{-0.31}^{+0.34}$ & 0.82$_{-0.02}^{+0.03}$ & 0.89$_{-0.08}^{+0.09}$ & 57.26$\pm0.61$ & 1086 (819) \\ 1708$-$400 & 0148690101 & 0.95$\pm0.01$ & 0.48$\pm0.01$ & 4.46$_{-0.10}^{+0.11}$ & 1.49$\pm0.04$ & 0.29$\pm0.02$ & 27.42$_{-0.13}^{+0.08}$ & 1566 (1232) \\ 1841$-$045 & 0013340101 & 1.86$_{-0.13}^{+0.14}$ & 0.52$\pm0.03$ & 3.79$_{-0.46}^{+0.57}$ & 1.99$_{-0.20}^{+0.25}$ & 0.21$_{-0.04}^{+0.05}$ & 17.39$\pm0.83$ & 408 (391) \\ 1841$-$045 & 0013340201 & 1.90$\pm0.11$ & 0.51$\pm0.02$ & 3.97$_{-0.40}^{+0.48}$ & 1.81$_{-0.12}^{+0.14}$ & 0.25$_{-0.03}^{+0.04}$ & 17.16$\pm0.51$ & 583 (641) \\ 2259$+$586 & 0038140101 & $0.59\pm0.01$ & $0.353_{-0.004}^{+0.005}$ & $4.88_{-0.13}^{+0.15}$ & $0.75\pm0.02$ & $0.50_{-0.04}^{+0.05}$ & $12.34\pm0.07$ & 1167 (888) \\ 2259$+$586 & 0155350301 & $0.54\pm0.01$ & $0.390\pm0.006$ & $5.01_{-0.14}^{+0.15}$ & $0.85\pm0.02$ & $0.67_{-0.04}^{+0.05}$ & $33.01\pm0.21$ & 1531 (1053) \\ 2259$+$586 & 0203550701 & $0.59\pm0.03$ & $0.356_{-0.015}^{+0.013}$ & $4.97_{-0.39}^{+0.47}$ & $0.82_{-0.07}^{+0.08}$ & $0.43_{-0.10}^{+0.13}$ & $13.77_{-0.44}^{+0.44}$ & 457 (458) \\ \hline \multicolumn{9}{@{}l@{}}{\hbox to 0pt{\parbox{180mm}{\footnotesize \footnotemark[$*$] Object name of the AXPs; CXOU\,J010043.1$-$721134 (0100$-$721), 4U\,0142$+$614 (0142$+$614), 1RXS\,J170849.0$-$400910 (1708$-$400), 1E\,1841$-$045 and 1E\,2259$+$586. \par\noindent \footnotemark[$\dagger$] XMM-Newton observation ID. \par\noindent \footnotemark[$\ddagger$] $N_{\mathrm{H}}$ denotes the column density with 90 \% confidence level errors. \par\noindent \footnotemark[$\S$] $kT_{\mathrm{LT}}$ and $kT_{\mathrm{HT}}$ denote the blackbody temperatures with 90 \% confidence level errors. \par\noindent \footnotemark[$\|$] $R_{\mathrm{LT}}$ and $R_{\mathrm{HT}}$ denote the emission radii with 90 \% confidence level errors. \par\noindent \footnotemark[$\#$] $F$ denotes a flux in the energy range 2-10 keV in units of $10^{-12}$ ergs cm$^{-2}$ s$^{-1}$ with 68 \% confidence level errors. }\hss}} \end{tabular} \end{center} \end{table*} \begin{table*} \caption{Spectral parameters of the quiescent emissions of the AXPs observed by {\it Chandra}.}\label{tab:axp_spc_cxo} \begin{center} \begin{tabular}{lllllllll} \hline\hline Object\footnotemark[$*$] & ObsID\footnotemark[$\dagger$] & $N_{\mathrm{H}}$\footnotemark[$\ddagger$] & $kT_{\mathrm{LT}}$\footnotemark[$\S$] & $R_{\mathrm{LT}}$\footnotemark[$\|$] & $kT_{\mathrm{HT}}$\footnotemark[$\S$] & $R_{\mathrm{HT}}$\footnotemark[$\|$] & $F$\footnotemark[$\#$] & $\chi^{2}$ (d.o.f.) \\ & & (10$^{22}$ cm$^{-2}$) & (keV) & (km) & (keV) & (km) & & \\ \hline 0100$-$721 & 1881 & 0.06$_{-0.05}^{+0.06}$ & 0.33$_{-0.04}^{+0.04}$ & 9.65$_{-1.74}^{+2.98}$ & 0.65$_{-0.11}^{+0.23}$ & 1.42$_{-0.85}^{+1.25}$ & 0.12$_{-0.05}^{+0.05}$ & 172 (150) \\ 0100$-$721 & 4616 & $<0.04$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & 0.39$\pm0.02$ & 6.84$_{-0.56}^{+1.06}$ & 0.09$\pm0.01$ & 63 (68) \\ 0100$-$721 & 4617 & $<0.41$ & 0.18$_{-0.06}^{+0.25}$ & 21.90$_{-20.07}^{+173.99}$ & 0.44$_{-0.05}^{+21.23}$ & 5.53$_{-5.52}^{+1.24}$ & 0.12$\pm0.04$ & 66 (64) \\ 0100$-$721 & 4618 & $<0.18$ & 0.28$\pm0.1$ & 10.88$_{-3.73}^{+15.89}$ & 0.51$_{-0.09}^{+1.04}$ & 3.02$_{-2.85}^{+2.67}$ & 0.10$\pm0.09$ & 50 (65) \\ 0100$-$721 & 4619 & $<0.04$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & 0.41$_{-0.2}^{+0.02}$ & 6.40$_{-0.51}^{+0.88}$ & 0.11$\pm0.01$ & 47 (67) \\ 0100$-$721 & 4620 & $<0.03$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & 0.40$_{-0.02}^{+0.01}$ & 6.67$_{-0.41}^{+0.79}$ & 0.10$\pm0.01$ & 70 (68) \\ 1647$-$455 & 6283 & 2.54$_{-0.69}^{+0.81}$ & 0.49$\pm0.06$ & 0.52$_{-0.18}^{+0.31}$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & 0.15$\pm0.04$ & 23 (21) \\ 1647$-$455 & 5411 & 1.44$_{-0.28}^{+0.32}$ & 0.58$\pm0.05$ & 0.26$_{-0.05}^{+0.07}$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & 0.13$\pm0.01$ & 54 (44) \\ \hline \multicolumn{9}{@{}l@{}}{\hbox to 0pt{\parbox{180mm}{\footnotesize \footnotemark[$*$] Object name of the AXPs; CXOU\,J010043.1$-$721134 (0100$-$721) and CXOU\,J164710.2$-$455216 (1647$-$455). \par\noindent \footnotemark[$\dagger$] {\it Chandra} observation ID. \par\noindent \footnotemark[$\ddagger$] $N_{\mathrm{H}}$ denotes the column density with 90 \% confidence level errors. \par\noindent \footnotemark[$\S$] $kT_{\mathrm{LT}}$ and $kT_{\mathrm{HT}}$ denote the blackbody temperatures with 90 \% confidence level errors. \par\noindent \footnotemark[$\|$] $R_{\mathrm{LT}}$ and $R_{\mathrm{HT}}$ denote the emission radii with 90 \% confidence level errors. \par\noindent \footnotemark[$\#$] $F$ denotes a flux in the energy range 2-10 keV in units of $10^{-12}$ ergs cm$^{-2}$ s$^{-1}$ with 68 \% confidence level errors. }\hss}} \end{tabular} \end{center} \end{table*} \begin{table*} \caption{Spectral parameters of the quiescent emission and the bursts of AXP\,CXOU\,J164710.2$-$455216 observed by {\it Swift}.}\label{tab:axp_j1647_spc_swift} \begin{center} \begin{tabular}{llllllll} \hline\hline SeqNum\footnotemark[$*$] & $N_{\mathrm{H}}$\footnotemark[$\dagger$] & $kT_{\mathrm{LT}}$\footnotemark[$\ddagger$] & $R_{\mathrm{LT}}$\footnotemark[$\S$] & $kT_{\mathrm{HT}}$\footnotemark[$\ddagger$] & $R_{\mathrm{HT}}$\footnotemark[$\S$] & $F$\footnotemark[$\|$] & $\chi^{2}$ (d.o.f.) \\ & (10$^{22}$ cm$^{-2}$) & (keV) & (km) & (keV) & (km) & & \\ \hline 00230341000\footnotemark[$\#$] & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & 8.92$_{-1.62}^{+1.34}$ & 1.21$_{-0.35}^{+0.59}$ & 0.35$\pm0.12$ & 8 (12) \\ 00030806001 & 1.72$_{-0.20}^{+0.31}$ & 0.63$_{-0.12}^{+0.07}$ & 2.40$_{-0.44}^{+0.94}$ & 2.08$_{-0.93}^{+38.74}$ & 0.14$_{-0.13}^{+0.36}$ & 27.94$\pm10.25$ & 253 (242) \\ 00030806002 & 1.76$_{-0.55}^{+0.70}$ & 0.73$_{-0.08}^{+0.08}$ & 1.79$_{-0.44}^{+0.7}$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & 20.01$\pm4.36$ & 16 (17) \\ 00030806003 & 1.95$_{-0.53}^{+0.86}$ & 0.42$_{-0.15}^{+0.26}$ & 3.03$_{-1.9}^{+9.21}$ & 0.81$_{-0.09}^{+0.35}$ & 1.04$_{-0.81}^{+0.43}$ & 13.98$\pm6.28$ & 101 (92) \\ 00030806004 & 1.79$_{-0.7}^{+0.91}$ & 0.51$_{-0.21}^{+0.28}$ & 2.20$_{-1.66}^{+4.56}$ & 0.93$_{-0.38}^{+3.06}$ & 0.68$_{-0.67}^{+0.5}$ & 13.79$\pm12.1$ & 91 (74) \\ 00030806006 & 0.99$_{-0.32}^{+0.37}$ & 0.75$\pm0.06$ & 1.24$_{-0.23}^{+0.31}$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & 11.99$\pm1.2$ & 31 (33) \\ 00030806007 & 1.65$_{-0.39}^{+0.49}$ & 0.67$_{-0.06}^{+0.07}$ & 1.52$_{-0.33}^{+0.48}$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & 9.85$\pm1.34$ & 27 (29) \\ 00030806008 & 1.78$_{-0.42}^{+0.55}$ & 0.60$\pm0.05$ & 2.02$_{-0.46}^{+0.69}$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & 9.18$\pm1.33$ & 73 (64) \\ 00030806009 & 2.09$_{-0.46}^{+0.57}$ & 0.62$\pm0.06$ & 1.69$_{-0.39}^{+0.58}$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & 6.86$\pm1.38$ & 64 (55) \\ 00030806010 & 1.51$_{-0.32}^{+0.4}$ & 0.71$\pm0.06$ & 1.31$_{-0.25}^{+0.34}$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & 9.48$\pm0.81$ & 47 (59) \\ 00030806011 & 1.35$_{-0.24}^{+0.28}$ & 0.67$\pm0.04$ & 1.38$_{-0.20}^{+0.26}$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & 7.93$\pm0.61$ & 90 (93) \\ 00030806012 & 1.30$_{-0.24}^{+0.29}$ & 0.67$\pm0.05$ & 1.22$_{-0.20}^{+0.25}$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & 6.60$\pm0.45$ & 66 (75) \\ 00030806013 & 1.79$_{-0.40}^{+0.49}$ & 0.63$\pm0.06$ & 1.57$_{-0.36}^{+0.54}$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & 7.00$\pm0.97$ & 63 (70) \\ 00030806014 & 1.57$_{-0.51}^{+0.72}$ & 0.63$\pm0.09$ & 1.26$_{-0.38}^{+0.68}$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & 4.64$\pm1.43$ & 22 (31) \\ 00030806015 & 1.80$_{-0.42}^{+0.53}$ & 0.61$_{-0.06}^{+0.07}$ & 1.38$_{-0.34}^{+0.53}$ & $\cdot\cdot\cdot$ & $\cdot\cdot\cdot$ & 4.77$\pm0.84$ & 49 (47) \\ \hline \multicolumn{8}{@{}l@{}}{\hbox to 0pt{\parbox{180mm}{\footnotesize \footnotemark[$*$] {\it Swift} sequence number. \par\noindent \footnotemark[$\dagger$] $N_{\mathrm{H}}$ denotes the column density with 90 \% confidence level errors. \par\noindent \footnotemark[$\ddagger$] $kT_{\mathrm{LT}}$ and $kT_{\mathrm{HT}}$ denote the blackbody temperatures with 90 \% confidence level errors. \par\noindent \footnotemark[$\S$] $R_{\mathrm{LT}}$ and $R_{\mathrm{HT}}$ denote the emission radii with 90 \% confidence level errors. \par\noindent \footnotemark[$\|$] $F$ denotes fluxes in the energy ranges 15-150 keV in units of $10^{-6}$ ergs cm$^{-2}$ s$^{-1}$ for the burst observation of 00230341000 and 2-10 keV in units of $10^{-12}$ ergs cm$^{-2}$ s$^{-1}$ for other observations with 68 \% confidence level errors. \par\noindent \footnotemark[$\#$] Results for the burst. }\hss}} \end{tabular} \end{center} \end{table*} \clearpage %%% % See the manual for the detail. %%% | 7 | 10 | 0710.3816 |
0710 | 0710.1778_arXiv.txt | Using Magellan/IMACS images covering a 1.2 x 1.2 sq. degree FOV with seeing of 0.4"-0.6", we have applied convolution techniques to analyse the light distribution of 364 confirmed globular cluster in the field of NGC 5128 and to obtain their structural parameters. Combining these parameters with existing Washington photometry from Harris et al. (2004), we are able to examine the size difference between metal-poor (blue) and metal-rich (red) globular clusters. For the first time, this can be addressed on a sample of confirmed clusters that extends to galactocentric distances about 8 times the effective radius, R$_{eff}$, of the galaxy. Within 1 R$_{eff}$, red clusters are about $30\%$ smaller on average than blue clusters, in agreement with the vast majority of extragalactic globular cluster systems studied. As the galactocentric distance increases, however, this difference becomes negligible. Thus, our results indicate that the difference in the clusters' effective radii, r$_e$, could be explained purely by projection effects, with red clusters being more centrally concentrated than blue ones and an intrinsic r$_e$--R$_{gc}$ dependence, like the one observed for the Galaxy. | \label{sec:intro} Since sizes and structural parameters of globular clusters (GCs) in different GC systems (GCSs) have first been obtained, it has become clear that some of these properties correlate with global properties of their host galaxies \citep[see for example][]{jordan05,brodie06}. The existence of the so called fundamental plane relation for an increasing number of studied GCSs seems to confirm that GCs populate a narrow region in this parameter space \citep{djorgovski95,mclaughlin00,mclaughlin05,barmby07}. However, there are puzzling trends that are still awaiting confirmation and need to be addressed using larger samples of GCs. It is necessary to study structural parameters of GCs and GC-like objects in different environments before definitive statements can be made regarding their formation. Among the structural parameters that can be studied, the effective (or half-light) radius is of particular importance. Models have shown that this quantity remains fairly constant throughout the entire GC lifetime \citep{spitzer72,aarseth98}, making it a good indicator of proto-GC sizes that are still observable today. A decade ago, HST observations unveiled a systematic size difference between red and blue GCs \citep{kundu98}. Since then, multiple studies have found that the blue GCs are between $17\%-30\%$ larger than their metal-rich counterparts in both spirals and early--type galaxies \citep{kundu99,puzia99,larsen01,larsen_fb01,kundu01,barmby02,jordan05}. However, most of these studies have made use of HST observations and examine only the innermost regions of the galaxy or small fields in regions at galactocentric distances greater than the galaxy's effective radius. According to \cite{larsen03}, the systematic size difference between red and blue GCs is caused merely by a projection effect. Since red (metal-rich) GCs are found to be more centrally concentrated than blue (metal-poor ones) in early type galaxies \citep[][among others]{cote01,dirsch03,woodley05}, the red GCs will appear to lie, on average, at a smaller galactocentric distance. The red clusters will on average be smaller than the blue clusters assuming that both types shares the same relation between the GC size and galactocentric distance. The relation r $\sim \sqrt{\rm{R}_{gc}}$ was first found in the Milky Way by \cite{vandenbergh91}. In this scenario, the difference between the cluster sizes should be most apparent at small galactocentric distance and should decrease strongly beyond 1 galaxy effective radius \citep{larsen03}. Alternatively, \cite{jordan04} suggests that this effect could be explained by an intrinsic difference between metal-rich and metal-poor GCs. Assuming half-mass radii that are independent of metallicity, effects of mass segregation combined with a metallicity-dependent stellar lifetime should lead to different sizes between the blue and red clusters. The brightest stars would be more massive and more centrally concentrated for the metal-rich GCs. This scenario should have little to no dependence on a cluster's distance from the center of its parent galaxy. In a recent study, \cite{spitler06} analysed the GCS of NGC 4594 (Sombrero, at a distance of $ \sim 9$ Mpc) using a six-image mosaic from HST/ACS. They confirm that within the inner 2 arcmin (2.2 R$_{eff}$), the metal-rich GCs are, on average, $17\%$ smaller than the metal-poor clusters. However, the size difference becomes negligible at $\sim 3$ arcmin, corresponding to $\sim 3.4$ R$_{eff}$, where R$_{eff} =0.89$ arcmin \citep{baggett98}. To further understand the sizes of red and blue clusters, we need a homogeneous survey of a GCS with the ability to eliminate contaminating sources, high resolution to measure structural parameters, and over a large range in galactocentric distance. NGC 5128 is the nearest giant elliptical galaxy, at a distance of 3.8 Mpc \citep{mclaughlin07}. Its GCs are thus easily resolvable with sub-arcsecond seeing \citep{harris06}. In this paper we present effective radius results for 337 GCs from the \cite{woodley07} catalog that are confirmed GCs by either radial velocity measurement from various studies \citep[see the references in ][]{woodley07} or are resolved by HST/ACS images \citep{harris06}. We also present the effective radii of 27 GCs newly confirmed through radial velocity measurements using the Baade 6.5-m telescope with the instrument LDSS-2 (data in preparation for publication). This list represents a clean sample of confirmed clusters. All of these also have ellipticities less than 0.4 and effective radii less than 8 pc, both of which are consistent with {\it normal} GC properties in NGC 5128. We find that only an additional $2.4\%$ of GCs from the \cite{woodley07} catalog have effective radii greater than the 8 pc boundary we have imposed here (to be discussed in detail in G\'omez \& Woodley, 2008, in preparation). Those few GCs are not considered here as our purpose is to establish the effective radius trends within the bulk of the GC population. | \label{sec:conclusions} Using a contaminant-free sample of 364 GCs in NGC 5128, confirmed with radial velocity measurements or by resolved HST images, we have measured effective radii using ISHAPE. Our results indicate that the blue or metal-poor clusters do not show any significant r$_{e}$--R$_{gc}$ relation. However, the red or metal-rich GCs do show a steep relation in which red clusters within 1 R$_{eff}$ of the galaxy's light are $30\%$ smaller than the blue clusters. Beyond this distance there is no indication for a size difference between the two metallicity populations. This finding in NGC 5128, not previously seen in any other early-type galaxy, supports the more tentative findings of the Sombrero galaxy's GCS \citep{spitler06}. Both studies support the idea that the size differences are most likely caused by projection effects \citep{larsen03} and not by intrinsic physical differences between the two subgroups. Acknowledgements: M.G. and K.A.W. thank Dean McLaughlin for use of HST structural parameters in advance of publication. M.G. thanks the Dept. of Physics and Astronomy at McMaster University and especially Bill and Gretchen Harris for their hospitality. K.A.W. thanks NSERC and Bill Harris for financial support, and also the Depto. de F{\'i}sica at the Universidad de Concepci{\'o}n, especially Doug Geisler, for their hospitality. We thank the anonymous referee for her/his valuable suggestions and comments. \clearpage | 7 | 10 | 0710.1778 |
0710 | 0710.3850_arXiv.txt | The first map of interstellar acetylene (C$_2$H$_2$) has been obtained with the infrared spectrograph onboard the {\it Spitzer Space Telescope}. A spectral line map of the $\nu_5$ vibration-rotation band at 13.7 $\mu$m carried out toward the star-forming region Cepheus A East, shows that the C$_2$H$_2$ emission peaks in a few localized clumps where gas-phase CO$_2$ emission was previously detected with {\it Spitzer}. The distribution of excitation temperatures derived from fits to the C$_2$H$_2$ line profiles ranges from 50 to 200 K, a range consistent with that derived for gaseous CO$_2$ suggesting that both molecules probe the same warm gas component. The C$_2$H$_2$ molecules are excited via radiative pumping by 13.7 $\mu$m continuum photons emanating from the HW2 protostellar region. We derive column densities ranging from a few $\times$ 10$^{13}$ to $\sim$ 7 $\times$ 10$^{14}$ cm$^{-2}$, corresponding to C$_2$H$_2$ abundances of 1 $\times$ 10$^{-9}$ to 4 $\times$ 10$^{-8}$ with respect to H$_2$. The spatial distribution of the C$_2$H$_2$ emission along with a roughly constant $N$(C$_2$H$_2$)/$N$(CO$_2$) strongly suggest an association with shock activity, most likely the result of the sputtering of acetylene in icy grain mantles. | Acetylene constitutes a key ingredient in the production of large complex hydrocarbon molecules in the dense interstellar medium (Herbst 1995). Because acetylene has no permanent dipole moment, it lacks a rotational spectrum that could be observed at radio wavelengths; observations of interstellar acetylene have therefore been limited to mid-infrared studies of rovibrational bands, carried out from ground-(e.g., Evans, Lacy \& Carr 1991; Carr et al. 1995) and space-based (e.g., Lahuis \& van Dishoeck 2000; Boonman et al. 2003; Lahuis et al. 2007) observatories. Acetylene has been detected in the gas-phase - either in absorption (e.g., Carr et al. 1995; Lahuis \& van Dishoeck 2000) or in emission (e.g. Boonman et al. 2003, this paper) - mostly toward young stellar objects. C$_2$H$_2$ can be used as a tracer of warm (100 K to 1000 K) molecular gas along often complicated sightlines. C$_2$H$_2$ abundance estimates, which were sometimes a few orders of magnitude higher than the predictions of cold gas-phase steady-state chemical models, have led to a better understanding of the role that warm gas-phase chemistry (e.g., Doty et al. 2002; Rodgers \& Charnley 2001) and/or grain mantle processing (e.g., Ruffle \& Herbst 2000) can play in star-forming regions, both locally and in extra-galactic objects (Lahuis et al. 2007). In this Letter, we present the first detection of the $\nu_5$ band of acetylene (C$_2$H$_2$) at 13.7 $\mu$m toward the star forming region Cepheus A East using the Infrared Spectrograph (IRS) onboard the {\it Spitzer Space Telescope}. This is the first map of C$_2$H$_2$ obtained toward any interstellar gas cloud. Section 2 describes the observations and data analysis. Sections 3 and 4 compare the spatial distribution of C$_2$H$_2$ to those of gaseous CO$_2$ and H$_2$ $S$(2), and discuss the C$_2$H$_2$-emitting gas in the context of shock chemistry and local outflow activity. The presence of C$_2$H$_2$ on interstellar dust grains will also be discussed in the context of cometary ices composition. | Acetylene was previously observed toward low-to-high mass star-forming regions either in absorption (e.g., Lahuis \& van Dishoeck 2000) or in emission (e.g., Boonman et al. 2003) using the {\it Infrared Space Observatory} ({\it ISO}). Excitation temperatures ranging from $\sim$ 10 to 900 K and abundances with respect to H$_2$ ranging from a few $\times$ 10$^{-8}$ to a few $\times$ 10$^{-7}$ were derived. Steady-state models of gas-phase chemistry in cold (10-50K) dense (10$^3$-10$^5$ cm$^{-3}$) molecular clouds predict abundances for acetylene between a few $\times$ 10$^{-10}$ and 1$\times$ 10$^{-8}$ with respect to H$_2$ depending on the role that neutral-neutral destruction reactions may play (Bettens, Lee \& Herbst 1995; Lee et al. 1996). Similar abundances are predicted by models of gas-grain chemistry in quiescent clouds with the highest values obtained only after 10$^6$ years (Ruffle \& Herbst 2000). While such models could account for the observed abundances in the cold gas, they were unable to reproduce the enhancements observed toward much warmer gas components in those objects. Mechanisms such as C$_2$H$_2$ ice sublimation from grain mantles and/or C$_2$H$_2$ enhancements via warm gas-phase chemistry were then invoked (e.g., Carr et al. 1995; Doty et al. 2002; Boonman et al. 2003). For the NE outflow region in Cepheus A East, we derive C$_2$H$_2$/H$_2$ abundance ratios in the range 1 $\times$ 10$^{-9}$ to 4 $\times$ 10$^{-8}$, for an assumed H$_2$ column density of 1.5 $\times$ 10$^{22}$ cm$^{-3}$ (G\'omez et al. 1999). These values are averages along the observed sight-lines. The highest abundances are localized around and to the south of, the NE position as well as around the positions of the HW5/6 sources (see Fig.~2). This is precisely where the interaction between the {\it northeast} outflow and the ambient molecular clouds occurs, and it is in these regions that we observe strong H$_2$ $S$(2) emission, a tracer of warm shocked gas. Thus the spatial variation of the C$_2$H$_2$ abundance again strongly suggests an association with shock activity, perhaps as a result of (1) production in the gas phase via high temperature reactions, or (2) grain mantle sputtering. Models for chemistry in hot cores (Rodgers \& Charnley 2001; Doty et al. 2002) indicate that enhanced abundances of C$_2$H$_2$ ($\sim$ few $\times$ 10$^{-8}$) are expected in warm regions with $T \ge$ 200 K. The good correlations between H$_2$ $S$(2), gaseous CO$_2$ and C$_2$H$_2$ indicate that such high temperatures were reached in the warm gas at the passage of the non-dissociative shock. However, the chemical pathways leading to the production of C$_2$H$_2$ are slow, and enhanced abundances only occur after $\sim$ 10$^4$ years, a time scale much greater than that expected for shock heating of the gas ($\sim$ 300 years; e.g. Kaufman \& Neufeld 1996). Hence, enhanced production of C$_2$H$_2$ by high temperature gas-phase chemistry is unlikely to be predominant in the observed region. Thus the correlations shown in Fig.~3 argue in favor of grain mantle sputtering over gas-phase production as the origin of the C$_2$H$_2$ in Cepheus A East. This is the same production mechanism that we favored for gaseous CO$_2$ (Sonnentrucker et al. 2006). While models for the production of C$_2$H$_2$ in shocks are not available to our knowledge, our results further suggest that both C$_2$H$_2$ and CO$_2$ are released into the gas phase under very similar physical conditions. The gaseous $\rm C_2H_2/CO_2$ ratio is roughly constant, with a mean value of 0.08 for all sight-lines where we detected acetylene at the 1.5 $\sigma$ level. If this value reflects the composition of the grain mantle, and given a $N$(CO$_2$)$_{ice}$/$N$(H$_2$O)$_{ice}$ ratio in this source of 0.22 (Sonnentrucker et al. 2007, ApJ in press), then the required $N$(C$_2$H$_2$)$_{ice}$/$N$(H$_2$O)$_{ice}$ ratio is 0.02. This value is at least a factor 4 larger than those derived toward other star-forming regions (0.1-0.5\%, Evans et al. 1991; Lahuis \& van Dishoeck 2000) and those predicted by theoretical models (0.1-0.5\%, Hasegawa \& Herbst 1993; Ruffle \& Herbst 2000), and at least a factor 2 larger than the gaseous C$_2$H$_2$/H$_2$O ratios obtained in observations of cometary comae (0.1-0.9\%, Brooke et al. 1996; Weaver et al. 1999). We speculate that these discrepancies might result from (1) the destruction of CO$_2$ by reaction with atomic hydrogen in shocks faster than $\sim$ 30 km s$^{-1}$ (predicted by Charnley \& Kaufman, 2000) and/or (2) a greater efficiency for sputtering of C$_2$H$_2$ in slow shocks.\footnote{Although both these effects would be strongly dependent upon the shock velocity, the relative constancy of the $\rm C_2H_2/CO_2$ would not necessarily require any fine tuning of the shock velocity. In reality, any sight-line typically samples an ensemble of shocks with a {\it range} of shock velocities, and the constancy of the $\rm C_2H_2/CO_2$ would simply indicate that the admixture of shock velocities varies little from one sight-line to another (e.g. Neufeld et al.\ 2006).} In either case, the gaseous C$_2$H$_2$/CO$_2$ ratios we observed may exceed the solid C$_2$H$_2$/CO$_2$ ratio, and the $N$(C$_2$H$_2$)$_{ice}$/$N$(H$_2$O)$_{ice}$ ratio could be less than 0.02. Unfortunately, direct measurements of acetylene ice are not possible, the weak features expected from solid C$_2$H$_2$ being blended with much stronger features of CO and H$_2$O (Boudin et al. 1998). However, further observations of gaseous C$_2$H$_2$ at a higher signal-to-noise ratio would be very valuable as a probe of any variations in the $\rm C_2H_2/CO_2$ ratio which might provide important clues to the shock physics. | 7 | 10 | 0710.3850 |
0710 | 0710.1064_arXiv.txt | In this paper we report on the gas-phase abundance of singly-ionized magnesium (Mg II) in 44 lines of sight, using data from the {\it Hubble Space Telescope} ({\it HST}). We measure Mg II column densities by analyzing medium- and high-resolution archival STIS spectra of the 1240 \AA{} doublet of Mg II. We find that Mg II depletion is correlated with many line of sight parameters (e.g.~$\fHmol$, $E_{B-V}$, $\ebvdist$, $A_V$, and $\avdist$) in addition to the well-known correlation with $\nHavg$. These parameters should be more directly related to dust content and thus have more physical significance with regard to the depletion of elements such as magnesium. We examine the significance of these additional correlations as compared to the known correlation between Mg II depletion and $\nHavg$. While none of the correlations are better predictors of Mg II depletion than $\nHavg$, some are statistically significant even assuming fixed $\nHavg$. We discuss the ranges over which these correlations are valid, their strength at fixed $\nHavg$, and physical interpretations. | \label{s:MgII_intro} Magnesium is both a relatively abundant element in the Galaxy and an important component in most interstellar dust models. Mg I has an ionization potential of only 7.65 eV, while Mg II has an ionization potential of 15.04 eV. In H I regions, the dominant form of gas-phase interstellar magnesium should be Mg II. In H II regions, magnesium should be found primarily in the form of Mg III. Gas-phase Mg I is rare, even in $\Hmol$ regions. As is the case for other elements such as silicon and iron, the average gas-phase abundance of magnesium is much smaller than the assumed overall cosmic abundance of magnesium, implying that the majority of interstellar magnesium is tied up in dust. Therefore, variations in the gas-phase magnesium abundance are only capable of having minor effects on grain composition. Nevertheless, observed variations in the gas-phase magnesium abundance may still shed light on the physical conditions of interstellar clouds. Two major {\it Copernicus} surveys that included measurements of Mg II abundances and depletions were \citet{Murray} and \citet{Jenkins1986}. Both of these studies confirmed that magnesium depletion increases with increased average hydrogen volume density, $\nHavg=\NHtot/r$ (where $r$ is the line-of-sight pathlength), in the line of sight. \citet{Jenkins1986} also found several correlations between the depletions of magnesium and other elements, which they cited as secondary to the correlation between those depletions and $\nHavg$. In both studies, the depletion of magnesium was not strongly correlated to other line of sight parameters, such as $\ebv$ and the magnitude of the 2175 \AA{} extinction ``bump''. We also note that there is a systematic difference in the absolute values of the abundances and depletions between these studies and more recent studies that is fairly substantial. This is due to a difference in the assumed $f$-values of the 1240 \AA{} doublet of magnesium. We comment on our chosen $f$-values in \S \ref{ss:MgII_lines}. More recently, \citet{Cartledge2006} also examined the abundances and depletions of Mg II and several other elements in the interstellar medium using STIS data. (Hereafter, this paper will be referred to as CLMS, for the initials of the authors.) Similar to the {\it Copernicus} studies, CLMS concluded that the line of sight parameter with the clearest connection to elemental depletions is the average hydrogen volume density, $\nHavg$. CLMS also explored potential correlations between depletions and other line of sight parameters such as the molecular fraction of hydrogen, $\fHmol$; selective extinction, $\ebv$; and selective extinction divided by line-of-sight pathlength, $\ebvdist$. CLMS concluded that $\nHavg$ was the parameter that best identified warm vs.~cold clouds, and that no other parameters produced correlations with magnesium depletion that were both as strong and with as little scatter. We began this study before CLMS was published. In spite of the similarities between that study and this one, we have proceeded with this study to provide an independent analysis of Mg II, and also because we have analyzed potential trends with respect to parameters not analyzed in CLMS, e.g.~$\av$ and $\rv$. The Mg II column densities of 11 out of the 44 lines of sight in our sample have also been analyzed previously by CLMS. We still report our results in this paper for two main reasons: (1) these lines of sight provide a basis of comparison for the methods of this paper and those of CLMS and (2) these lines of sight can be analyzed with respect to the aforementioned line of sight parameters not analyzed in CLMS. A comparison of the column density measurements for these common lines of sight is found in \S \ref{ss:MgII_coldensities}. In \S \ref{s:MgII_obsdata} we discuss our observations and data reduction, including comments on the 1240 \AA{} doublet and our derivation of column densities and abundances. In \S \ref{s:results} we discuss our results, including observed correlations and a review of the Galactic abundance of magnesium. In \S \ref{s:MgII_summary} we summarize our findings. | \label{s:results} \subsection{Correlations} \label{ss:correlations} Our results for the column densities of Mg II are given in Table \ref{MgII_coldensities}. The major correlation found by both \citet{Jenkins1986} and CLMS is that the depletion of magnesium and many other elements increases with overall line of sight density, $\nHavg$. Both interpreted their models in light of the model by \citet{Spitzer1985}. The \citeauthor{Spitzer1985} model contends that the ISM contains two distinct varieties of cloud (warm and cold), each with a distinct depletion level. In this model, lines of sight with average densities of a few-tenths of a particle per cm$^{-3}$ are largely sampling the warm ISM, while lines of sight with average densities of a few particles per cm$^{-3}$ are largely sampling the cold ISM. This explains the observed plateaus of depletion at both low and high densities, with a transition near average densities of $\sim1$ cm$^{-3}$, where lines of sight are sampling both types of clouds. This model, however, does not take into account the transition to a much different regime of cloud chemistry expected for translucent clouds \citep[for a recent review, see][]{SnowMcCall}. \citet{Jenkins1986} also found evidence of lesser correlations between the depletion of various elements and other parameters of reddening and extinction, but concluded that these correlations were secondary to the correlation with $\nHavg$. CLMS examined the depletion of Mg II compared to $\fHmol$, $\ebv$, and $\ebvdist$. CLMS concluded that only $\ebvdist$ provided a correlation with magnesium depletion that was as at least roughly as significant as the correlation with $\nHavg$, but with increased scatter. However, in our recent work \citep{JensenFeII}, we explored the possibility of correlations between iron depletion and various other line of sight parameters. We found very similar correlations between iron depletion and measures of dust density ($\ebvdist$ and $\avdist$) and the molecular fraction of hydrogen ($\fHmol$). We also examined iron depletion as a function of $\rv$, the ratio of total visual extinction to selective extinction, though a conclusive trend did not present itself. Using many of the same lines of sight in this paper as in \citet{JensenFeII}, we find the same correlations generally hold with respect to magnesium depletion. In what follows, we discuss the nature of these correlations and possible interpretations. \subsubsection{Correlations with Hydrogen} \label{sss:hydrogen_correlations} First, we note that magnesium depletion is clearly correlated with total hydrogen volume density $\nHavg$. Five lines of sight (HD 27778, HD 37021, HD 37061, HD 37903, and HD 147888) have substantially larger depletions than any of the other lines of sight in this sample; HD 27778, HD 37021, HD 37061, and HD 147888 are the four densest lines of sight (in terms of $\nHavg$), while the density of HD 37903 is in the top 20\% of our sample. This correlation is plotted in Figure \ref{fig:logMgIIHlognh}. Since these dense lines of sight are also some of the lines of sight that are found in both our sample and that of CLMS, we are not probing significantly higher average density and cannot expand on their conclusions regarding $\nHavg$. We also find that magnesium depletion is correlated with total hydrogen column density. However, we conclude that this is primarily a secondary correlation due to the correlation with $\nHavg$. However, we can attempt to analyze whether or not the correlation is independently significant. To do this, we follow the methods described in \citet{Jenkins1986}. First, we calculate Pearson correlation coefficients between depletion and the two variables of interest (in this case $\logHtot$ and $\lognHavg$), as well as those two variables with each other. The partial correlation coefficient is then given by $\rho_{12.3} = (\rho_{12}-\rho_{13}\rho_{23})/[(1-\rho_{13}^2)(1-\rho_{23}^2)]^{-1/2}$, where the subscripts on the correlation coefficients indicate the two variables being correlated, and $\rho_{12.3}$ is the correlation coefficient between the first two variables if the third variable is held fixed. However, what is really being calculated is $r$, the sample correlation coefficient(s), as opposed to $\rho$, the population correlation coefficient(s). Once $r_{12.3}$ has been calculated, we examine the significance level for a $t$-test of the appropriate number of degrees of freedom (in this case, the number of data points minus three) to determine the probability that $\rho$ is non-zero (i.e.~a true correlation exists). Using these methods, we find a 49\% chance of the null hypothesis (i.e.~the two-sided probability that $\rho=0$) for a correlation between Mg II depletion and $\NHtot$ when $\nHavg$ is held fixed. Because this technique of trivariate analysis uses Pearson correlation coefficients, a few caveats apply, namely that linear relationships and normal distributions are implicitly assumed. We used $\logMgIIH$, $\logHtot$, and $\lognHavg$ in the analysis just discussed because the correlation between all combinations of those variables are stronger than when in linear form. CLMS also explored possible correlations between magnesium depletion and $\fHmol$, but only briefly commented on the results. They concluded that any correlation was less significant than the correlation between magnesium depletion and $\nHavg$, in that it did not as effectively discriminate between the distinct depletion levels expected in the \citet{Spitzer1985} model. The upper right panel of Figure 9 in CLMS shows a very clear correlation between magnesium depletion and $\fHmol$ for $\fHmol \gtrsim 0.1$, superimposed with a scatter plot in depletion for a smaller subset of lines with $0.01 \lesssim \fHmol \lesssim 0.1$. Two additional lines of sight with $\fHmol \lesssim 10^{-4}$ conform to the main correlation in that they show minimal depletion. We also see a clear correlation between magnesium depletion and the molecular fraction of hydrogen, $\fHmol$, plotted in Figure \ref{fig:logMgIIHHf}, with depletion increasing with increasing $\fHmol$. There is the exception of one discrepant point, HD 147888, that exhibits substantial depletion at $\fHmol \sim 0.1$. We note that CLMS found a larger column density for this line of sight than we do ($0.2\dex$); however, even if we adopt the CLMS Mg II column density for this line of sight, its abundance is still $0.26\dex$ smaller than any line of sight in our sample with $\fHmol < 0.4$. We concur with CLMS that the correlation between magnesium depletion and $\fHmol$ is not as rigorous as the correlation with $\nHavg$. We examine the data using various combinations of the variables in their logarithmic (where correlations are the strongest) and linear forms. In all cases, we find that the probability of the null hypothesis (i.e.~$\rho=0$, as discussed above) between Mg II depletion and $\fHmol$ with $\nHavg$ held constant is less than 5\%; in most cases, it is $\lesssim1$\%. Conversely, the probability that there is no correlation between Mg II depletion and $\nHavg$ with $\fHmol$ held constant is less than 0.1\%. Therefore, we conclude that while $\nHavg$ is clearly the dominant correlation, the correlation with $\fHmol$ has some independent significance. A question of interest is where scatter seems to be introduced and why. Between Figure 9 of CLMS and Figure \ref{fig:logMgIIHHf} of this paper, scatter is only observed at $\fHmol \lesssim 0.1$. In our recent related work on Fe II \citep{JensenFeII}, we note that Fe II depletion is also correlated with $\fHmol$, but there are a few exceptions to the trend at both high and low values of $\fHmol$. The most severe exceptions noted in that paper were HD 147888 and HD 164740 (with large depletions but $\fHmol \lesssim 0.1$) and HD 210121 (with less depletion despite $\fHmol \sim 0.7$). For Mg II, however, we do not see any outlying points at values of $\fHmol$ larger than $\sim0.1$ within either this sample or CLMS. Points such as HD 147888, however, still require explanation. \citet{Snow1983} put forth the possibility that in some dense environments an increased average grain size, which decreases the grain surface area per unit volume, may suppress $\Hmol$ formation (as $\Hmol$ is thought to form on grain surfaces). This scenario was specifically discussed in the context of the $\rho$ Oph cloud, for which there is indepedent evidence (in part, a value of $\rv$ greater than the interstellar average of 3.1) that grain coagulation has occurred. Our main outlying line of sight, HD 147888 (with $\rv$ of 4.06), passes through the $\rho$ Oph cloud, so observing the combination of a dense, depleted environment with small $\fHmol$ is not surprising in this case. Two of the other lines of sight with large depletions are HD 37021 and HD 37061. Reliable far-ultraviolet data sets do not exist for these lines of sight; therefore, they do not have measurements of the molecular hydrogen column densities or subsequently derived molecular fractions of hydrogen. However, as stated above in \S \ref{ss:hydrogen}, \citet{Cartledge2001} has argued that these lines of sight have small values of $\fHmol$ based on a lack of Cl I. These two lines of sight also have large values of $\rv$ (5.54 and 4.23, respectively), implying a larger average grain size. The possible effect of grain size on depletions and $\Hmol$ formation, and other interpretations, will be discussed further in \S \ref{sss:extinction_correlations}. Barring such outlying points, the trend of increased Mg II depletion with increased $\fHmol$ has a relatively clear interpretation. $\Hmol$ is formed in the same dense, dusty environments that foster large depletions. Within the context of this sample and CLMS, this seems to hold for lines of sight with $\fHmol \gtrsim 0.1$. While the ubiquity of $\Hmol$ even in diffuse regions complicates the issue \citep[see conclusions of][]{Rachford2002}, there is still a physical argument that $\fHmol$ should be a good diagnostic of the local conditions of interstellar clouds, and therefore depletions. It is worth noting that the similar trend between iron depletion and $\fHmol$ exhibits scatter up to $\fHmol \sim 0.3$ in the work of \citet{SavageBohlin}, in addition to the outlying points mentioned above from \citet{JensenFeII}. Whether the range in $\fHmol$ over which there is scatter in the abundances is truly different for magnesium and iron or is simply a selection effect is unclear. \subsubsection{Correlations with Extinction and Reddening Parameters} \label{sss:extinction_correlations} We find that the depletion of Mg II is correlated to both selective extinction, $\ebv$, and total visual extinction, $\av$. However, there is significant scatter in these correlations. Because these are integrated line of sight parameters, it makes sense to divide by line-of-sight pathlength. Both $\ebv$ and $\av$ are strongly correlated with $\NHtot$, and both are thought to be rough measures of the total dust column density. Therefore, $\ebvdist$ and $\avdist$ should be strongly correlated with $\nHavg$ and be approximations of the total dust volume density. When we look for correlations between magnesium depletion and $\ebvdist$ and $\avdist$ we find that the correlations are substantially increased when compared to the integrated line of sight parameters. Therefore, we can say, with reasonable confidence, that magnesium depletion is increased in increasingly dusty environments. The correlations with $\ebvdist$ and $\avdist$ are plotted in Figures \ref{fig:logMgIIHebv_dist} and \ref{fig:logMgIIHav_dist}. As with the correlation between depletion and $\fHmol$, we examine partial correlation coefficients to determine the independent significance of these correlations. The partial correlation coefficient between Mg II depletion and $\log{\ebvdist}$ with $\lognHavg$ held fixed implies that the probability of the null hypothesis is less than 6\%. The same partial correlation coefficient with $\log{\avdist}$ implies that there is less than a 1\% chance of the null hypothesis. (Note that both probabilities are two-sided to a $t$-distribution.) We consider these variables in their logarithmic forms because these are the versions of the variables that exhibit the strong correlations (for all combinations of the variables in question). However, if the variables are considered in non-logarithmic forms, the probability of the null hypothesis generally increases. Again we note that this type of trivariate statistical measure implicitly assumes that the correlation is linear with normally distributed scatter. As with the correlation between depletion and $\fHmol$, we conclude that while these correlations do not improve on $\nHavg$ as a predictor of depletions, this is limited evidence that they are significant in their own right. We have briefly explored the possibility that the additional (or ``missing'') magnesium that is depleted from these lines of sight is found in the form of gas-phase Mg I in regions that are presumably shielded from radiation by dust. While the STIS data do not cover Mg I absorption lines in many cases, our results indicate that gas-phase Mg I column densities are far too small to account for the order-of-magnitude increase in depletion seen in gas-phase Mg II. Therefore, it likely that the missing gas-phase magnesium is tied up in the additional grains found in these environments, supporting the conclusion above that the correlations between Mg II depletion and the parameters $\log{\ebvdist}$ and $\log{\avdist}$ are physically significant. It is also worth noting that in no case do we see Mg/H less than $\sim1\ppm$, even in the densest environments, with values of $\ebvdist$ and $\avdist$ several times larger than the average of the sample. This suggests that these lines of sight are not probing what might be considered ``translucent clouds'' \citep[though some may be ``translucent lines of sight''; see][]{SnowMcCall}. The Mg II abundance is plotted against the ratio of total visual to selective extinction, $\rv \equiv \av / \ebv$, in Figure \ref{fig:logMgIIHRv}. The plot shows significant scatter; however, some statistical measures show the possibility of a slight correlation. A Pearson correlation coefficient between $\logMgIIH$ and $\rv$ implies that the probability of the null hypothesis is about 31\%. A Spearman's $\rho$ rank correlation coefficient, which does not depend on the functional form assumed (including whether variables are considered linearly or logarithmically) beyond assuming that the correlation is either monotonically increasing or monotonically decreasing, shows a negative correlation (decreasing abundance/increasing depletion as $\rv$ increases) and is significant to approximately 1.3-$\sigma$. Therefore, there is evidence of a possible slight correlation between depletion and $\rv$. However, this is far from a certain conclusion. Significant selection effects are also a possibility, as the correlations are dominated by some of the points that have greater depletion and large values of $\rv$. In fact, if these lines of sight are excluded, the trend begins to reverse toward a positive correlation between increasing Mg II abundance and increasing $\rv$. In general, we conclude that $\rv$ is a poor predictor of depletions; as one anecdotal counterexample, HD 91597 has a very large value of $\rv=4.9$ but does not exhibit particularly large Mg II depletion. However, there are a few lines of sight that present interesting interpretive challenges where the value of $\rv$ may provide insight. As discussed above, the possibility of large magnesium depletion but small $\fHmol$ exists for three lines of sight: HD 37021, HD 37061, and HD 147888. In the latter case the effect is clear, while in the former two cases the small value of $\fHmol$ is merely inferred. What is interesting, as noted above, is that these three lines of sight all have values of $\rv>4$ which is a fairly significant deviation from the interstellar average of 3.1. Because large grains contribute to $\av$ (i.e.~grey extinction) but less so to $\ebv$, $\rv$ is thought to be correlated to average grain size. Explaining why depletion should increase in a line of sight with large grains is difficult. As mentioned above in our discussion of $\Hmol$ and iron depletion, increased grain size decreases dust surface area per unit volume, and therefore reduces $\Hmol$ formation rates. However, decreased surface area per unit volume also implies a reduction in rates of sticking between dust grains and gas-phase atoms. It seems we can reasonably conclude that the large values of $\rv$, i.e.~the larger average grain populations, are not responsible for the large depletions by way of atoms and ions sticking to the grains. One possibility is that the large depletions are instead ``locked in'' prior to grain coagulation. Another possibility is the effect of a high-radiation field: this is known for the line of sight toward HD 147888 ($\rho$ Oph D) as well as HD 37021 and HD 37061 which are in Orion (radiation is presumed to be responsible for the relative lack of Cl I, and thus also $\Hmol$, as mentioned in \S \ref{ss:hydrogen}). However, \citet{Snow1983} argues that the increased radiation is unlikely to be entirely responsible for the low $\fHmol$ in the $\rho$ Oph cloud. Whether or not this is the case for HD 37021 and HD 37061 is unclear. More details about the radiation field and the exact nature of the grain population (we have only considered the crude measure of $\rv$) are probably necessary to fully understand these lines of sight. \subsubsection{Anticorrelation with Distance} \label{sss:distance_anticorrelation} We find that magnesium depletion is generally anticorrelated with distance to the background star, that is, line of sight pathlength; this relationship is shown in Figure \ref{fig:logMgIIHdist}. Depletion decreases by nearly an order of magnitude between very short lines of sight and those up to about 2 kpc or so, and then is relatively constant (to within about 0.3-0.4$\dex$) out to about 6 kpc. As we concluded for a similar anticorrelation seen between iron depletion and distance \citep{JensenFeII}, the long pathlengths are likely sampling a variety of cloud conditions, resulting in the constant depletion for long-pathlength lines of sight. On the other hand, given the comparable hydrogen column densities of all the lines of sight in this study ($\logHtot \approx21-22$), the shorter lines of sight are generally the denser lines of sight. \subsubsection{Spatial Variations} \label{sss:spatial_variations} We find one very interesting correlation with Galactic location: the five stars with the largest depletions reside at higher Galactic latitudes of $|b|>15^{\circ}$. However, with pathlengths of less than 1 kpc, these lines of sight are still primarily in the Galactic disk. When we analyze magnesium depletions against the height from the center of the Galactic disk, $z=r \sin{b}$, we do not see a strong correlation. The variation with respect to Galactic latitude is most likely a coincidence, given that these are some of the densest and most reddened lines of sight. We do not see any other evidence of significant spatial variations. \subsection{Mg/H of Galactic Stars} \label{ss:stellarMgH} In the last several years, three major papers have attempted to analyze the cosmic abundance ``standards'' in the ISM through studies of stellar abundances and meteoritic abundances---\citet{SnowWitt}, \citet{SofiaMeyer}, and \citet{Lodders}. The importance of these standards is to compare them with the observed gas-phase abundances and infer an absolute value for depletions---and therefore absolute values for the amount of these elements in phases other than atomic gas, i.e.~dust grains and molecules. Of the major elements relevant to dust, the element with the best determined cosmic abundance is iron. The four major measurements of the cosmic Fe/H ratio---solar, B stars, F and G stars, and CI chondrites---all agree very closely, largely within the errors. The situation is somewhat more complex for other elements. Carbon and oxygen show apparent overabundances in the Sun compared to F and G stars (whether or not the solar and F/G star abundances potentially agree within the errors depends on the choice of solar abundances, regarding which there is still some uncertainty), while B stars show relative deficits in these abundances compared to the Sun and other F and G stars. The chondritic abundances of C and O are even smaller. Nitrogen seems to be somewhat less abundant in B stars than in the Sun, though the two are reconciliable within the errors; F and G nitrogen abundances are generally unknown, and the chondritic abundances are substantially lower. Silicon seems to be most abundant in F and G stars, slightly less abundant in the Sun, and about half as abundant in B stars. The errors, however, do not rule out agreement between all three measurements. However, the chondritic abundance tightly matches the solar abundance. Both \citet{SofiaMeyer} and \citet{Lodders}, cite \citet{Holweger} for the solar abundance of magnesium, $\logMgH=-4.46$; \citet{SnowWitt} report a slightly older value from \citet{AndersGrevesse} of $\logMgH=-4.42$, though these values are consistent within the errors. The chondritic abundances in \citet{Lodders} of $\logMgH=-4.44$ are also very consistent with these solar values. The discrepancy arises when various stellar abundances are considered. Both \citet{SnowWitt} and \citet{SofiaMeyer} found significantly smaller abundances of Mg for B stars. \citet{SnowWitt} found $\logMgH=-4.63$ for field B stars and $\logMgH=-4.68$ for cluster B stars; \citet{SofiaMeyer}, making no distinction between cluster and field stars, found $\logMgH=-4.64$ for all B stars. Though there is marginal agreement within the very large errors in these numbers, the B star abundances are $\approx60\%$ smaller than the solar abundances. \citet{SnowWitt} and \citet{SofiaMeyer} also disagree on the Mg abundance in F and G stars ($\logMgH=-4.52$ and $-4.37$, respectively) due most likely to \citet{SofiaMeyer} restricting their sample to stars with ages of $\leq2$ Gyr. Again, however, these values have relatively large errors and are reconciliable with the B star abundances, though just barely. What effect does the choice of a cosmic magnesium abundance have for the implied dust-phase abundances to be used in dust models? The differences between the cosmic abundances just discussed leads to nontrivial differences in the dust-phase abundances. Our weighted interstellar average of Mg II/H is $2.7\pm0.1\ppm$ (parts per million), though the median value in our 44 lines of sight is somewhat larger at $6.2\ppm$. Taking the extremes of the above numbers, anywhere between $\sim20$ and $\sim40\ppm$ of Mg is available for creating dust. Examining the various models compared in Table 3 of \citet{SnowWitt}, we find that most models require much more Mg than the lower value of $\sim20$ implied by a B star abundance standard. That a B star abundance is less likely to represent the cosmic abundance was also found by \citet{ZDA2004}, who had more difficulty fitting dust models to observations using the dust-phase abundances implied by assuming B star abundances as the cosmic standard. In fact, this is true even though \citeauthor{ZDA2004} assumed $\approx0$ ppm of magnesium to be in the gas-phase. If the few ppm of magnesium in the gas-phase as measured by this paper and CLMS were included, the \citet{ZDA2004} fits would become even more strained (the best fits for B star abundances were at the limit of the error in those abundances and inferior to the fits obtained using other abundances). Therefore, the major conclusion that we can make regarding cosmic abundances and the incorporation of magnesium into dust is to add to the evidence that B star abundances, despite B stars being younger and therefore potentially good tracers of the current ISM, are a poor cosmic standard. Whether or not the solar or an F and G star abundance standard for Mg is a better fit is a test that is too sensitive for us to comment on, given the uncertainties in those abundances. We have analyzed the abundance of Mg II in 44 lines of sight. Our study does not probe substantially larger average hydrogen volume densities than previously observed by CLMS; therefore, we observe the same correlations between Mg II and the $\nHavg$ and $\ebvdist$. We also note a correlations between magnesium depletion and $\avdist$, a different measure of dust density. Correlations between $\NHtot$ and the reddening and extinction parameters $\ebv$ and $\av$ mean that correlations between Mg II depletion and dust density measures are expected. However, these latter correlations, while not strong than the correlation between depletion and $\nHavg$, show some evidence of being significant even at fixed $\nHavg$ and should be more directly related to the line-of-sight dust content. We also note a correlation between magnesium depletion and $\fHmol$ in our data; combined with the results of CLMS, this correlation seems to be valid for $\fHmol \gtrsim 0.1$ but not at smaller $\fHmol$. A question that is related to the trend with $\fHmol$ is why so little $\Hmol$ forms in certain high-density lines of sight. Our results are consistent with the \citet{Snow1983} suggestion that the reduced grain surface area per unit volume of large grains plays a role in reducing $\Hmol$ formation rates. For similar reasons, we can conclude that the grain coagulation probably occurs after depletions are already ``locked into'' the dust, rather than depletion of gas-phase atoms onto grain surfaces. | 7 | 10 | 0710.1064 |
0710 | 0710.1913_arXiv.txt | We report on the H$_2$O maser distributions around IRAS 22480+6002 (=IRC+60370) observed with the Japanese VLBI Network (JVN) at three epochs spanning 2 months. This object was identified as a K-type supergiant in 1970s, which was unusual as a stellar maser source. The spectrum of H$_2$O masers consists of 5 peaks separated roughly equally by a few km s$^{-1}$ each. The H$_2$O masers were spatially resolved into more than 15 features, which spread about 50 mas along the east--west direction. However, no correlation was found between the proper motion vectors and their spatial distributions; the velocity field of the envelope seems random. A statistical parallax method applied to the observed proper-motion data set gives a distance of $1.0\pm 0.4$ kpc for this object, that is considerably smaller than previously thought. The distance indicates that this is an evolved star with $L\sim 5800\ L_{\odot}$. This star shows radio, infrared, and optical characteristics quite similar to those of the population II post-AGB stars such as RV Tau variables. | \label{sec:introduction} H$_2$O maser emission has been observed in circumstellar envelopes of evolved stars such as O-rich Mira variables and OH/IR stars with large mass loss rates of $\dot{M} \geq 10^{-7}M_{\odot}$~yr$^{-1}$ \citep{rei81,eli92}. Most of these stars are asymptotic giant branch (AGB) stars or red supergiants both with the spectral type M, with a few exceptions for transient stars at pre-planetary nebula phase (or supposedly a few pre-main sequence stars such as Ori KL; \cite{mor98}). For the central stars with spectral types earlier than M, UV radiation from stellar chromosphere eventually dissociates most of molecules (except CO) in the inner envelope (e.g., \cite{wir98}). Therefore, H$_2$O (or SiO) masers are usually not expected for these stars, except for the case that the molecules in dense circumstellar clumps shield themselves from UV radiation. In fact, H$_2$O masers found in a young planetary nebula \citep{mir01,sue07} must be such an exceptional case. OH masers have been found in yellow hypergiants with spectral types F and G (such as IRC+10420 and V1427 Aql) \citep{gig76,ned92,hum02}. However, H$_2$O and SiO masers have never been detected in these objects \citep{nak03}, though thermal emission of a few other molecules have been observed in the outer circumstellar shell \citep{cas01,tey06}. \begin{table*}[ht] \caption{Status of the telescopes, data reduction, and resulting performances in the individual epochs of the JVN observations.}\label{tab:status} \begin{center} \footnotesize \begin{tabular}{lccccccc} \hline \hline & Epoch in & & & Reference & 1-$\sigma$ level & Synthesized & Number of \\ Observation & the year & Duration & Used & velocity\footnotemark[2] & noise & beam\footnotemark[3] & detected \\ code & 2005 & (hr) & telescopes\footnotemark[1] & (km s$^{-1}$) & (Jy beam$^{-1}$) & (mas) & features \\ \hline r05084b \dotfill & March 25 & 7.3 & MZ, IR, OG, IS, KS, NB\footnotemark[4] & $-52.3$ & 0.22 & 1.7$\times$1.6, $-$37$^{\circ}$ & 20 \\ r05116b \dotfill & April 26 & 7.3 & MZ, IR, OG, IS\footnotemark[5], KS, NB & $-52.0$ & 0.15 & 3.8$\times$2.0, $-$14$^{\circ}$ & 17 \\ r05151a \dotfill & May 31 & 8.1 & MZ, OG\footnotemark[5], IS\footnotemark[5], KS, NB & $-$52.6 & 0.15 & 3.2$\times$2.8, $-$66$^{\circ}$ & 14 \\ \hline \end{tabular} \end{center} \footnotemark[1] Telescopes that were effectively operated and whose recorded data were valid: MZ: the VERA 20~m telescope at Mizusawa, IR: the VERA 20~m telescope at Iriki, OG: the VERA 20~m telescope at Ogasawara Is., IS: the VERA 20~m telescope at Ishigakijima Is., KS: the NiCT 34-m telescope at Kashima, NB: the NRO 45-m telescope at Nobeyama. \\ \footnotemark[2] Velocity channel used for the phase reference in data reduction. \\ \footnotemark[3] The synthesized beam made in natural weight; major and minor axis lengths and position angle. \\ \footnotemark[4] Ceasing operation for 2.5~hr due to strong winds and pointing correction. \\ \footnotemark[5] High system temperature ($>$300~K) due to bad weather conditions. \end{table*} \ \\ The optical counterpart of IRAS 22480+6002 ($=$AFGL~2968, or IRC$+$60370) was identified as a K-type supergiant (K0Ia; \cite{hum74}, or K4.5Ia; \cite{faw77}). Therefore, the detections of H$_2$O and SiO masers \citep{han98,nym98} were surprising. Though a search for OH 1612 MHz emission was negative \citep{les92}, CO emission was detected toward this star \citep{jos98}. From the CO $J=2$--1 line profile, the systemic stellar velocity and the expansion velocity of this star were obtained to be $V_{\rm lsr}$=$-49.3$ km s$^{-1}$ and $V_{exp}=26.4$ km~s$^{-1}$, respectively \citep{jos98, gro99}. They are consistent with those obtained from the H$_2$O\ and SiO maser spectra, and the expansion velocity of the envelope of this star is typical for OH/IR stars. The radial velocity gives a kinematic distance of 5.0~kpc. It suggests a large luminosity $L_{\ast}=$140 000~$L_{\odot}$ of the central star \citep{gro99}, but it is consistent with the supergiant interpretation of this object. A blue nearby star, a B5II star, is seen by about 12$''$ east of this object. Though it is cataloged as a visual binary \citep{wor97}, a physical association of this object with the maser source is questionable because of the large velocity difference of about 40 km s$^{-1}$ \citep{hum74}. \citet{win94} gave a new spectral classification of M0I for IRAS 22480+6002 from the low-resolution spectrum between 6000 and 8800 A, which was significantly different from the previous type assignment of this star. For a long-period variable, optical spectral classification may vary with light variations. However, this star has not been reported as a variable star, though it is optically not very faint ($V\sim 8.3$). In this work, we report three-epoch VLBI observations of H$_2$O\ masers of IRAS 22480+6002 to rectify the entangled situation associated with this object. From the spatio-kinematics of the masers, we diagnose a probable anomaly of a hot wind from the K-type star. We estimated the distance to this star using the statistical parallax method based on the proper motion data of H$_2$O masers. Our result gives a much smaller distance for this star than previously thought. The new estimation of the distance demands to reconsider various properties of this star. Based on the arguments presented in section 3, we conclude that this star is a population II post-AGB star. | \begin{table*}[ht] \caption{Parameters of the H$_2$O maser features identified by proper motion toward IRAS 22480$+$6002.} \label{tab:pmotions} \begin{center} \begin{tabular}{lrrrrrrrrrrr} \hline \hline Feature\footnotemark[1] & \multicolumn{2}{c}{Offset} & \multicolumn{4}{c}{Proper motion\footnotemark[2]} & \multicolumn{2}{c}{Radial motion\footnotemark[3]} & \multicolumn{ 3}{c}{Peak intensity at 3 epochs} \\ & \multicolumn{2}{c}{(mas)} & \multicolumn{4}{c}{(mas yr$^{-1}$)} & \multicolumn{2}{c}{(km s$^{-1}$)} & \multicolumn{ 3}{c}{(Jy beam$^{-1}$)} \\ & \multicolumn{2}{c}{\hrulefill} & \multicolumn{4}{c}{\hrulefill} & \multicolumn{2}{c}{\hrulefill} & \multicolumn{ 3}{c}{\hrulefill} \\ & $\Delta$R.A. & $\Delta$decl. & $\mu_{x}$ & $\sigma(\mu_{x})$ & $\mu_{y}$ & $\sigma(\mu_{y})$ & $V_{z}$ & $\Delta V_{z}$\footnotemark[4] & Epoch 1& Epoch 2& Epoch 3 \\ \hline 1 \ \dotfill \ &$ -6.56$&$ -10.69$&$ 4.20$& 1.46 &$ 5.86$& 1.40 &$ -59.78$& 0.26 & 0.19 & 0.17 & ... \\ 2 \ \dotfill \ &$ 0.00$&$ 0.00$&$ 0.00$& 0.36 &$ 0.00$& 1.29 &$ -58.48$& 2.74 & 2.33 & 2.72 & 3.48 \\ 3 \ \dotfill \ &$ -45.23$&$ 7.26$&$ -0.87$& 0.42 &$ 0.51$& 0.66 &$ -55.01$& 2.04 & 3.77 & 2.47 & 2.42 \\ 4 \ \dotfill \ &$ -4.95$&$ 7.17$&$ 0.25$& 0.29 &$ 0.47$& 1.44 &$ -54.33$& 1.05 & 0.41 & 0.53 & 0.73 \\ 5 \ \dotfill \ &$ -26.70$&$ -2.86$&$ -1.03$& 0.97 &$ -0.72$& 0.61 &$ -52.34$& 3.37 & 8.11 & 8.73 & 8.16 \\ 6 \ \dotfill \ &$ -0.37$&$ 1.50$&$ 3.96$& 1.29 &$ 0.86$& 1.91 &$ -50.37$& 0.56 & 0.53 & 0.73 & 0.49 \\ 7 \ \dotfill \ &$ 0.16$&$ 2.74$&$ 0.05$& 0.61 &$ 0.98$& 1.56 &$ -49.29$& 2.18 & 8.44 & 7.96 & 6.20 \\ 8 \ \dotfill \ &$ -0.04$&$ 13.96$&$ 0.90$& 0.67 &$ -0.29$& 0.78 &$ -47.88$& 1.33 & 1.92 & 1.69 & 1.03 \\ 9 \ \dotfill \ &$ -34.92$&$ -7.63$&$ -1.95$& 1.66 &$ -3.07$& 2.97 &$ -47.39$& 0.52 & 0.25 & 0.26 & ... \\ 10 \ \dotfill \ &$ -48.23$&$ 11.08$&$ 2.27$& 1.18 &$ 0.71$& 0.88 &$ -47.24$& 0.56 & 0.24 & 0.20 & 0.18 \\ 11 \ \dotfill \ &$ -2.24$&$ 8.18$&$ 0.06$& 1.10 &$ -0.16$& 0.64 &$ -46.17$& 0.84 & 0.41 & 0.28 & 0.26 \\ 12 \ \dotfill \ &$ -1.02$&$ 13.29$&$ 0.06$& 0.41 &$ 0.56$& 0.74 &$ -45.37$& 1.97 & 3.34 & 3.40 & 2.44 \\ 13 \ \dotfill \ &$ 0.13$&$ 7.77$&$ -1.86$& 3.19 &$ 0.26$& 1.32 &$ -44.40$& 0.77 & 0.24 & 0.33 & 0.32 \\ \hline \end{tabular} \end{center} \noindent \footnotemark[1] H$_2$O maser features detected toward IRAS 22480+6002. The feature is designated as IRAS 22480+6002:I2007 {\it N}, where {\it N} is the ordinal source number given in this column (I2007 stands for sources found by Imai et~al. and listed in 2007). \\ \footnotemark[2] Relative value with respect to the motion of the position-reference maser feature: IRAS 22480+6002:I2007 {\it 2}. \\ \footnotemark[3] Relative value with respect to the local stand of rest. \\ \footnotemark[4] Mean full velocity width of a maser feature at half intensity. \end{table*} \subsection{Spatial distribution and proper motions of maser features} Figure \ref{fig:I2248_spectrum} shows cross-power spectra of the H$_2$O masers of IRAS 22480+6002. The H$_2$O maser emission spread in a velocity range of 15 km~s$^{-1}$, which is typical for Mira-type AGB stars (e.g., \cite{tak94}). Five spectral peaks were seen in roughly equal separations of 2--3 km~s$^{-1}$; the second highest peak was near the systemic velocity ($V_{\rm lsr}=-49.3$ km s$^{-1}$). The correlated powers of these peaks equally increased by a factor of two during 2 months in our observing run, except for the second peak for which the intensity increased only by about 20\%. However, the peak flux densities of individual features were found not to vary much (Table 2). This fact indicates that extended emissions were partially resolved in the shortest baseline between NRO and NICT (197.4 km), but resolved-out in the longer baselines. % Correlated flux densities were estimated to be about 30--40\% of the total-power intensities. Figure \ref{fig:1st-epoch} shows the distribution of maser features at the first epoch. The extent of 50 mas corresponds to 50~AU ($D$/1 kpc), which is somewhat larger than those seen around Mira variables at $D=1$ kpc (e.g., \cite{bow94}). In this figure, one of the low-velocity (blue-shifted) components ($V_{\rm lsr}=-58.48$ km s$^{-1}$: the position reference) is located at the origin and the other low velocity components (in grey and blue colors) are located both near the eastern and western edges. Many of the higher velocity (red-shifted) components (shown in yellow, orange, and red colors) fall at the eastern edge, but a few of them are scattered at the west side too. The overall distribution of water maser features are characterized by the elongation to the east-west direction. But, no clear correlation is found between the velocities and spatial positions. If the circumstellar envelope of the K supergiant interacts with the wind from the eastern BII star (though this is unlikely), maser spots and features could be aligned perpendicularly to the wind direction, i.e., in the north--south direction (e.g., \cite{ima02b}). We find no such N--S alignment of the H$_2$O maser features. \citet{mei99} noted that the MIR image of this star with the NASA 3-m telescope showed a northeast-southwest elongation, but concluded that it was likely to be an artifact caused by astigmatism. \begin{figure}[htb] \begin{center} \FigureFile(8cm,8cm){fig2.eps} \end{center} \caption{Distribution of H$_2$O masers on 2005 March 25. The color code indicates the radial velocity of the feature, and the size of the filled circle indicates the flux density of the feature. Note that the $-58.48$ km s$^{-1}$ reference component is located at the origin (light blue), but is almost overlapped with the systemic-velocity component ($\sim -49$ km s$^{-1}$) shown in green.} \label{fig:1st-epoch} \end{figure} We detected 14--20 H$_2$O maser features though all epochs (the last column of table \ref{tab:status}). Note that the H$_2$O masers were persistent in velocity and in spatial distribution during the two months; 65--90\% of the detected maser features survived during our observing run. Therefore, we identified the same maser features at three epochs relatively easily, and measured the proper motions of the individual features during two months. Table \ref{tab:pmotions} gives the measured proper motions. Figure \ref{fig:PM-I2248} shows the linear fits to the relative positions of the individual maser features. The fitted proper motions look significant with the second epoch contributing little for most features, which warrants our selecting the same maser features at the different epochs. % Circumstellar H$_2$O masers can amplify the radiation of the central star (for example, see the case of U Her; \cite{vle02}). % In the present case, we may speculate that the $-52.34$ km s$^{-1}$ feature (No. 5 in Table 2 and Figure 4), which is one of the strongest components and located near the center of the maser distribution, is such a maser amplifying the steller radiation. However, it is hasty to draw any conclusion from this observation, since we have no information on the central-star position in this scale. \begin{figure}[ht] \FigureFile(8cm,14cm){fig3.eps} \caption{Observed relative proper motions of H$_2$O maser features in IRAS 22480$+$6002 in R.A. (a) and decl. (b) directions. The number on the left indicates the feature name designated in Table 2. The vertical bar attached to each data point indicates the position uncertainty. The least-square-fitted line is also shown.} \label{fig:PM-I2248} \end{figure} \begin{figure}[ht] \FigureFile(8cm,6cm){fig4.eps} \caption{Doppler velocities (colorfully displayed) and relative proper motion vectors (indicated by arrows) of H$_2$O masers in IRAS 22480$+$6002. The displayed proper motion vector is that subtracted by a velocity bias ($\overline{\mu _{x}}, \overline{\mu _{y}})=(0.97, 0.72)$ [mas yr$^{-1}$] from the original vector to cancel out the average motions of all the features. A number added for each feature with a proper motion shows the assigned one after the designated name form ``IRAS 22480$+$6002: I2007". The map origin is set to the location of the feature IRAS 22480$+$6002: I2007 {\it 2}.} \label{fig:I2248-velocity} \end{figure} Figure \ref{fig:I2248-velocity} shows the proper motion vectors of the individual H$_2$O maser features. Note that the largest two proper-motion vectors at the lower left and lower right, i.e., features 1 and 9 of the $V_{\rm lsr}=-59.78$ and $-47.39$ km s$^{-1}$ components, respectively, were determined by two-epoch detections, so that they are slightly inaccurate. The proper motions of all other features with 3-epoch detections are within a few mas per year (relative to the reference component at $V_{\rm lsr}=-58.48$ km s$^{-1}$). We cannot find any systematic trend of motions in this diagram. For example, features 3 and 10 at the western edge move in opposite directions, and features 6 and 13 at the eastern edge also move in opposite directions. Figure \ref{fig:expansion} shows the RA-offsets and $\mu_x$ plots against $V_{\rm lsr}$. The ellipse is a plot of expected offset and proper motion from a thin spherical-shell model with a constant velocity (in the Right Ascension direction because the maser features are spread mainly in this direction). If the shell model is correct, all of the maser features should fall between these ellipses. However, the right panel does not show such a tendency. Rather, the observed points seem to distribute randomly. The randomness of the proper motions may partially originate from the large random errors in the position measurements. % In order to check this issue, we made a Monte Carlo simulation of the 3-epoch proper motion fitting with the same positional uncertainties but without real motions (i.e., position jitters only due to the measurement errors). We obtained the mean velocity dispersions ($0.79\pm 0.25$ mas yr$^{-1}$, $0.80\pm 0.20$ mas yr$^{-1}$) for the 13 proper motions in R.A. and Dec. directions for the present case from the simulations; here, the number after the '$\pm$' sign is a standard deviation of dispersions obtained in the simulations. The observed velocity dispersions (1.95 mas yr$^{-1}$, 1.93 mas yr$^{-1}$), are significantly larger than the simulated mean dispersions (more than $4 \sigma$). Therefore, the observed proper motions are substantially real motions of masing features. \begin{figure}[ht] \FigureFile(8cm,6cm){fig5.eps} \caption{Plot of relative R.A. offset (left) and proper motion (right) against radial velocity. The filled and unfilled circles indicate the three-epoch and two-epoch detections. The large and small ellipses in both panels indicate the position and proper motion curves expected from thin spherical shell models (one-dimensional in the R.A. direction) with a constant expansion velocity of 12.5 km s$^{-1}$ and a radius of $3.5\times 10^{14}$ cm, and with 7 km s$^{-1}$ and $8\times 10^{13}$ cm, respectively both at a distance of 0.9 kpc. The observed points in the right panel do not fit to these ellipses. } \label{fig:expansion} \end{figure} The observed random motions may originate from the intrinsic random ballistic motions of matters ejected from or infalling into the atmosphere of the central supergiant. This may be a characteristic of mass outflow of a supergiant originating from the extended atmosphere which is considerably turbulent \citep{lev05,jos07}. The motion of 1 mas yr$^{-1}$ corresponds a transverse motion of $\sim 4.7 (D/$kpc) km s$^{-1}$. In order to obtain the distance, we applied the statistical parallax method to the obtained maser proper motions; % for example, see \citet{sch81}. % Assuming random motions of maser features, we obtained a velocity dispersion in radial motion (with respect to the average velocity of maser features, $V_{\rm lsr}=-50.6$ km s$^{-1}$) to be $\sigma_v \simeq 5.0$ km s$^{-1}$, which is smaller than the outflow velocity estimated from CO emission. The dispersion in the maser proper motions can be obtained by subtracting the dispersion involved in the measurements; % see equation (3) of \citet{sch81}. % We obtain $\sigma _{\mu} \sim 1.40$ mas yr$^{-1}$ and get a distance to IRAS 22480+6002 to be $D=\sigma_{v}/\sigma_{\mu}= 0.76 \ (\pm 0.25)$ kpc. The formal uncertainty involved in the distance estimation was computed using equation (4) of \citet{sch81}. If we exclude the largest two proper-motion features with two-epoch detections from the sample, we get the distance $1.02 \ (\pm 0.38) $ kpc. Later on, we adopt this distance for IRAS 22480+6000, because large motions detected by two-epoch observations are somewhat dubious. Note that this distance is derived based on the assumption that the velocity field of masers is random and isotropic, and that the proper motions appeared in maser features are real motions of gas clumps. The distance 1.0 kpc gives a radius of water maser shell approximately $3.7\times 10^{14}$ cm, which is compatible with the radii of water maser shells of miras, but considerably smaller than those of M-supergiants \citep{yat94,cot04}. Though the obtained distance still involve a considerable uncertainty, it excludes the possibility of a very large distance of 5 kpc (a kinematic distance). The luminosity of this star is re-evaluated to be $5.8 \times 10^3~L_{\odot}$ (reestimated from \cite{gro99}). It is considerably small for a supergiant. However, from the distance of 1.0 kpc, we can compute the absolute V magnitude of this star from 2MASS K magnitude ($K=2.8$) using $V-K=3.7$ (for K5III; \cite{zom90}), and with reddening corrections, we get $M_V\sim -4.4$. This value falls near the absolute magnitude of K5Ib (or M0Ib) \citep{zom90}. Therefore, the luminosity is still in a range of supergiants. The radial velocity of $\sim -50$ km s$^{-1}$ is typical for young objects in the Perseus spiral arm in the direction of this star (for example, see \cite{sit03}). If we take into account the large uncertainties involved in the obtained distance, we cannot completely exclude the possibility that IRAS 22480+6002 belongs to the Perseus arm at $D\sim 3.0$ kpc at $l=108^{\circ}$ \citep{xux06}. However, there are several other bright stellar maser sources with similar radial velocities in the same direction, e.g., CU Cep ($-50$ km s$^{-1}$), IRC+60374 ($-52$ km s$^{-1}$), and AFGL 2999 ($-50$ km s$^{-1}$). Luminosity distances to these stars are inferred to be smaller than 3 kpc from their high IRAS flux densities. In addition, MY Cep is an M supergiant with $V_{\rm lsr}=-56$ km s$^{-1}$ in the star cluster NGC 7419. The distance to this cluster has been well estimated to be about 2.3 kpc from luminosities of member stars of the cluster (e.g., see \cite{bea94,sub06}). These example indicates that the stars with $V_{\rm lsr}\sim -50$ km s$^{-1}$ do not necessarily belong to the Perseus arm, but they may be located much closely. Because the radial velocity expected by the galactic rotation is only $\sim -10$ km s$^{-1}$ at 1 kpc at $l=108^{\circ}$, and because the radial-velocity dispersion of stellar maser sources is as small as $\sim 25$ km s$^{-1}$ at the solar neighborhood (see Appendix 2 of \cite{deg05}), IRAS 22480+6002 is possibly kinematically anomalous. \begin{table*}[ht] \begin{center} \caption{Comparison of the catalogued positions of IRAS 22480$+$6002.} \label{tab:positions} \footnotesize \begin{tabular}{llllllll} \hline \hline Catalog & Band & Assignment & epoch & R.A.(J2000) & decl..(J2000) & error & flux density \\ & & & & \ h \ m \ s & \ \ ${\circ} \ \ ' \ \ ''$ & $''$ & or magnitude \\ \hline IRAS & MIR & 22480+6002 & 1983 & 22 49 59.2 & +60 17 55 & $11''\times 5''$ (19$^{\circ}$) & $F_{\rm 12}=142$ Jy \\ MSX6 & MIR & G108.4255+00.8939 & 1995 & 22 49 58.89 & +60 17 56.8 & 0.3 & $F_{\rm C}=123$ Jy\\ & & & & & & & \\ 2MASS & NIR & 2495897+6017567 & 1997 & 22 49 58.97 & +60 17 56.8 & 0.29 & K=2.78 \\ GSC1.2 & optical& 0426500695 & 1954 & 22 49 59.43 & +60 17 55.8 & 0.3 & R=12.29 \\ GSC2.2 & optical& N012302336407 & 1989.6 & 22 49 58.900 & +60 17 57.17 & 0.3 & B=12.29 \\ & & & & & & & \\ USNO-B1.0 & optical& 1502-0356025 & 1971.7 & 22 49 59.75 & +60 17 56.7 & (0.7, 1.0) & R=8.87 \\ USNO-B1.0 & optical& 1502-0356023 & 1979.7 & 22 49 59.44 & +60 17 55.9 & (0.7, 0.2) & R=8.73 \\ USNO-B1.0 & optical& 1502-0356019 & 1979.7 & 22 49 59.15 & +60 17 57.5 & (0.5, 0.7) & B=12.63 \\ & & & & & & & \\ JVN (this work) & radio & 22480+6002 & 2005.5 & 22 49 58.876 & +60 17 56.65 & $0.1''$ & $F_{\rm H_2O}\sim 8$ Jy \\ \hline \hline \end{tabular} \end{center} \end{table*} \subsection{Past optical/infrared data of IRAS 22480+6002 (=J22495897+6017568).} Though this star is relatively bright at optical wavelengths ($V\sim 8.30$; Tyco Input catalog), the star was not recorded in major optical catalogs, for example, not in Henry Draper (HD) Catalogue, The Hipparcos and Tycho Catalogue, nor the General Catalog of Variable Stars, possibly because of confusion by the nearby B5II star (TYC 4265-870-1; $V\sim 10.74$), located by about 12$''$ east. This was involved in The Washington Double Star Catalog\footnote{available at http://ad.usno.navy.mil/wds/wdstext.html.}, giving a separation of 10.9$''$ in 1901 and 12.0$''$ in 2006 with a small position angle variation (by $\sim 7^{\circ}$) to the B star. From this data, we obtain the proper motion of 17 mas yr$^{-1}$ to the west for IRAS 22480+6002 relative to this B star. As noted by \citet{hum74}, this B5II star is probably not a binary counterpart because of the large radial velocity difference. The ACT Reference Catalog gave a very small proper motion of this B5II star (less than 3 mas yr$^{-1}$); though The Hipparcos and Tycho Catalogue gave a large proper motion in declination due to position uncertainty, but this was corrected in ACT catalog. We also checked the past catalogs recording the position of this star and summarized the results in table 3.\footnote{ all the data except JVN are available in the VizieR database ({\it http://vizier.nao.ac.jp/viz-bin/VizieR}).} The GSC 1.2 catalog (which remeasured the POSS1 plate taken in 1950s) gave a different position by about 4.2$''$, which leads a large proper motion of 84 mas yr$^{-1}$ if compared with GSC 2.0. This value is much larger than the above-mentioned proper motion computed from the Washington Double Star Catalog, though the proper motion vectors are roughly in the same direction. We believe the direct measurements of binary separation gives better values. Therefore, we adopt the proper motion of 17 mas yr$^{-1}$ for this star, and get $U_0=-71$ km s$^{-1}$ and $V_0=-29$ km s$^{-1}$ for IRAS 22480+6002. This motion is considerably peculiar for a population I disk star. It is likely that IRAS 22480+6002 belongs to one of kinematical streaming groups of stars as population II G and K giants \citep{fam05}. In the past, OH and H$_2$O masers have been found in a few planetary and preplanetary nebulae \citep{zil89,mir01,sue07}, where central stars of these objects have spectral types earlier than M. These masers are a remnant of circumstellar material which was ejected at the AGB phase of the central star. The molecules responsible for masers are eventually to be dissociated. In contrast, \citet{fix84} found OH 1665/1667 MHz emission toward several warm stars as RV Tau variables with spectral type of G and K, but so far only one case (TW Aql, a semi-regular variable of K7III) was confirmed to be a circumstellar maser \citep{pla91}. SiO masers, which are emitted within a few stellar radii of the central star (much closer than H$_2$O masers are emitted), were not detected in these warm objects before. An exceptional case is the RV Tau variable, R Sct, with spectral type K0Ib. This object exhibits strong SiO and weak H$_2$O masers (I. Yamamura, 2004 private communication) as well as the 4 $\mu$m SiO first overtone bands \citep{mat02}. The RV Tau variables are believed to be low-mass post-AGB stars \citep{jul86} with low metal abundances (population II; \cite{gir00}), though these are spectroscopically classified as supergiants. Their spectral types change between K and M-type with light variation \citep{pol97}. The atmosphere of late-type supergiants are not in hydrostatic equilibrium; effective temperature increases with decreasing metalicity \citep{lev05}. The RV Tau variables are enshrouded by dust shell, and CO emission has been detected in two of these variables \citep{buj88}. Although the optical counterpart of IRAS 22480+6002 is not reported to show any strong light variability (e.g., TASS; The Amateur Sky Survey\footnote{data available at http://www.tass-survey.org/}), the optical spectroscopic classification, middle infrared properties, and maser characteristics of IRAS 22480+6002 indicate a close similarity to the properties of the RV Tau variables. | 7 | 10 | 0710.1913 |
0710 | 0710.4525_arXiv.txt | We report a sensitive search for the \bhcn\ emission line towards SDSS\,J114816.64+525150.3 (hereafter:\ J1148+5251) at $z$=6.42 with the Very Large Array (VLA). HCN emission is a star formation indicator, tracing dense molecular hydrogen gas ($n({\rm H_2}) \geq 10^4\,$cm$^{-3}$) within star-forming molecular clouds. No emission was detected in the deep interferometer maps of J1148+5251. We derive a limit for the HCN line luminosity of $L'_{\rm HCN} < 3.3 \times 10^{9}\,$K \kms pc$^2$, corresponding to a HCN/CO luminosity ratio of $L'_{\rm HCN}$/$L'_{\rm CO}$$<$0.13. This limit is consistent with a fraction of dense molecular gas in J1148+5251 within the range of nearby ultraluminous infrared galaxies (ULIRGs; median value:\ $L'_{\rm HCN}$/$L'_{\rm CO}$=0.17$^{+0.05}_{-0.08}$) and HCN-detected $z$$>$2 galaxies (0.17$^{+0.09}_{-0.08}$). The relationship between $L'_{\rm HCN}$ and $L_{\rm FIR}$ is considered to be a measure for the efficiency at which stars form out of dense gas. In the nearby universe, these quantities show a linear correlation, and thus, a practically constant average ratio. In J1148+5251, we find $L_{\rm FIR}$/$L'_{\rm HCN}$$>$6600. This is significantly higher than the average ratios for normal nearby spiral galaxies ($L_{\rm FIR}$/$L'_{\rm HCN}$=580$^{+510}_{-270}$) and ULIRGs (740$^{+505}_{-50}$), but consistent with a rising trend as indicated by other $z$$>$2 galaxies (predominantly quasars; 1525$^{+1300}_{-475}$). It is unlikely that this rising trend can be accounted for by a contribution of active galactic nucleus (AGN) heating to $L_{\rm FIR}$ alone, and may hint at a higher median gas density and/or elevated star-formation efficiency toward the more luminous high-redshift systems. There is marginal evidence that the $L_{\rm FIR}$/$L'_{\rm HCN}$ ratio in J1148+5251 may even exceed the rising trend set by other $z$$>$2 galaxies; however, only future facilities with very large collecting areas such as the Square Kilometre Array (SKA) will offer the sensitivity required to further investigate this question. | High redshift galaxy populations are now being detected back to 780 million years after the Big Bang (spectroscopically confirmed:\ $z$=6.96; Iye \etal\ \citeyear{iye06}), probing into the epoch of cosmic reionization (e.g., Fan \etal\ \citeyear{fan06}; Hu \& Cowie \citeyear{hu06}). Many of these very distant galaxies show evidence for star formation activity (e.g., Taniguchi \etal\ \citeyear{tan05}). Some are even found to be hyperluminous infrared galaxies (HLIRGs; Bertoldi \etal\ \citeyear{ber03a}; Wang \etal\ \citeyear{wan07}) with far-infrared (FIR) luminosities exceeding 10$^{13}$\,\lsol, suggesting vigorous star formation and/or AGN activity. To probe the earliest stages of galaxy formation and the importance of AGN in this process, it is necessary to study the star formation characteristics of these galaxies. A good diagnostic to examine the star-forming environments in distant HLIRGs are observations of molecular gas, the fuel for star formation. The by far brightest and most common indicator of molecular gas in galaxies is line emission from the rotational transitions of carbon monoxide (CO), which was detected in $\sim$40 galaxies at high redshift ($z$$>$1; see Solomon \& Vanden Bout \citeyear{sv05} for a review). These observations have revealed molecular gas reservoirs with masses of $>$10$^{10}$\,\msol\ in these galaxies, even in the highest redshift quasar known, J1148+5251 at $z$=6.42 (Walter \etal\ \citeyear{wal03}, \citeyear{wal04}; Bertoldi \etal\ \citeyear{ber03b}). Although CO is a good tracer of the total amount of molecular gas in a galaxy, due to the relatively low critical density of $n_{\rm H_2} \sim 10^2-10^3\,$cm$^{-3}$ required to collisionally excite its lower $J$ transitions, it is not a reliable tracer of the dense molecular cloud cores where the actual star formation takes place. Recent studies of nearby actively star-forming galaxies have shown that hydrogen cyanide (HCN) is a far better tracer of the dense ($n_{\rm H_2} \sim 10^5-10^6$\,cm$^{-3}$) molecular gas where stars actually form (e.g.\ Gao \& Solomon \citeyear{gao04a}, \citeyear{gao04b}, hereafter:\ GS04a, GS04b). In the local universe it was found that the HCN luminosity ($L'_{\rm HCN}$) scales linearly (unlike $L'_{\rm CO}$) with the FIR luminosity ($L_{\rm FIR}$) over 7--8 orders of magnitude, ranging from Galactic dense cores to ULIRGs (Wu et al.\ \citeyear{wu05}). As $L_{\rm FIR}$ traces the massive star formation rate (unless AGN heating is significant), this implies that HCN is also a good tracer of star formation. HCN has now also been detected in five galaxies at $z$$>$2 (Solomon \etal\ \citeyear{sol03}; Vanden Bout \etal\ \citeyear{vdb04}; Carilli \etal\ \citeyear{car05}, hereafter:\ C05; Wagg \etal\ \citeyear{wag05}; Gao \etal\ \citeyear{gao07}, hereafter:\ G07). Adding a number of upper limits obtained for other high-$z$ galaxies, these observations indicate that the more luminous, higher redshift systems systematically deviate from the linear $L'_{\rm HCN}$--$L_{\rm FIR}$ correlation found in the local universe (G07), and hint at a rising slope of the relation toward high $L_{\rm FIR}$ and/or $z$. To further investigate this apparent non-linear, rising trend, our aim has been to extend the range of existing HCN observations beyond redshift 6 and to higher $L_{\rm FIR}$. In this letter, we report sensitive VLA\footnote{The Very Large Array is a facility of the National Radio Astronomy Observatory, operated by Associated Universities, Inc., under cooperative agreement with the National Science Foundation.} observations of \bhcn\ emission toward the $z$=6.42 quasar J1148+5251, the highest redshift source detected in CO. A previous, less sensitive search for \bhcn\ emission in this source has yielded no detection (C05). We use a concordance, flat $\Lambda$CDM cosmology throughout, with $H_0$=71\,\kms\,Mpc$^{-1}$, $\Omega_{\rm M}$=0.27, and $\Omega_{\Lambda}$=0.73 (Spergel \etal\ \citeyear{spe06}). | \subsection{Median Gas Density and Star Formation Efficiency} Krumholz \& Thompson (\citeyear{kt07}) argue that $L_{\rm FIR}$/$L'_{\rm HCN}$ is expected to be higher for galaxies with a median (molecular) gas density $n_{\rm med}$ close to or higher than the critical density $n_{\rm crit}^{\rm HCN}$ required for excitation of the observed HCN transition than for galaxies with lower $n_{\rm med}$. In their case, they define star formation efficiency as the fraction of the mass that is converted into stars per dynamical time of the system. Note that this is different than the star formation rate per unit total gas mass. They argue that the non-linear relation between $L_{\rm FIR}$ and $L'_{\rm CO}$ (e.g., Kennicutt \citeyear{ken98a}, \citeyear{ken98b}; GS04b; Riechers \etal\ \citeyear{rie06b}) arises due to the fact that CO traces all gas. The star-formation rate is then dictated by the density $n$ divided by the free-fall time $\tau_{\rm ff}$ ($\tau_{\rm ff} \propto n^{-0.5}$), giving the standard Schmidt-law: star formation rate $\propto n^{1.5}$, or $L_{\rm FIR}$ $\propto (L'_{\rm CO})$$^{1.5}$. For molecules like HCN, which only trace the small fraction of dense gas clouds directly associated with star formation in normal galaxies, $\tau_{\rm ff}$ is roughly fixed by $n_{\rm crit}$. Hence the star formation rate shows a linear relationship with $n$, or $L_{\rm FIR}$ $\propto (L'_{\rm HCN})$$^{1.0}$. However, in extreme galaxies, where $n_{\rm med}$ in the molecular ISM approaches $n_{\rm crit}^{\rm HCN}$, $\tau_{\rm ff}$ again becomes relevant (i.e., HCN emission no longer selects just the rare, dense peaks whose density is fixed by $n_{\rm crit}^{\rm HCN}$, but instead traces the bulk of the ISM, whose density can vary from galaxy to galaxy, and thus the variation of $n$ and $\tau_{\rm ff}$ re-enter the calculation), and the relationship approaches $L_{\rm FIR}$ $\propto (L'_{\rm HCN})$$^{1.5}$ (and $L'_{\rm HCN}$ $\propto L'_{\rm CO}$). Interestingly, current data show a marginal trend for a changing power-law index at the highest luminosities of the type proposed by Krumholz \& Thompson. This change in power-law index from 1 to 1.5 would suggest that, in these extreme luminosity systems, $n_{\rm med}$ approaches $n_{\rm crit}^{\rm HCN}$. More systems at high luminosity are required to confirm this trend of changing power-law index. \subsection{The Role of AGN Heating for $L_{\rm FIR}$} Like most of the $z$$>$2 HCN-detected sources, J1148+5251 is a quasar. It has been found that, even for such strong AGN galaxies, the bulk of $L_{\rm FIR}$ is likely dominatly heated by star formation in most cases (e.g., Carilli \etal\ \citeyear{car01}; Omont \etal\ \citeyear{omo01}; Beelen \etal\ \citeyear{bee06}; Riechers \etal\ \citeyear{rie06b}). However, based on radiative transfer models of the dust SED of J1148+5251, Li \etal\ (\citeyear{li07}) argue that this source may currently undergo a `quasar phase', in which AGN heating of the hot and warm dust contributes significantly to $L_{\rm FIR}$. If correct, this may be an alternative explanation for the elevated $L_{\rm FIR}$/$L'_{\rm HCN}$ in this galaxy. The (rest-frame) IR properties (tracing emission from hot dust) of J1148+5251 are similar to those of other $z$$>$6 quasars with much lower $L_{\rm FIR}$ (tracing emission from warm dust), and even to local quasars (Jiang \etal\ \citeyear{jia06}). This supports the assumption that the hot dust in J1148+5251 is dominantly heated by the AGN; however, the lack of a correlation between $L_{\rm IR}$ and $L_{\rm FIR}$ in quasars indicates that the warm dust may still be dominantly heated by star formation. Moreover, J1148+5251 follows the radio-FIR correlation for star-forming galaxies (Carilli \etal\ \citeyear{car04}), which also suggests a starburst origin for the dominant fraction of $L_{\rm FIR}$. Furthermore, one of the $z$$>$2 HCN detections and some of the meaningful limits are submillimeter galaxies without a known luminous AGN, but are still offset from the local $L_{\rm FIR}$/$L'_{\rm HCN}$ relation. It thus appears unlikely that AGN heating alone can account for the higher average $L_{\rm FIR}$/$L'_{\rm HCN}$ in the high-$z$ galaxy sample. \subsection{Implications for Future Studies} Even when assuming the highest $L_{\rm FIR}$/$L'_{\rm HCN}$ of 2835 found among all HCN-detected galaxies in Table \ref{tab-2}, the depth of our observations is sufficient to detect a galaxy with the redshift and $L_{\rm FIR}$ of J1148+5251 ($z$=6.42) in HCN emission at a signal-to-noise ratio of $>$4.5. To first order, our lower limit thus is consistent with previous suggestions (G07) that $L_{\rm FIR}$/$L'_{\rm HCN}$ ratios in high redshift sources lie systematically above those for nearby galaxies. The scatter around this trend is still significant, and will primarily be improved by increasing the number of HCN-detected galaxies at high $z$. In addition, it will be important to improve on the main sources of error for the individual high-$z$ detections (e.g., signal-to-noise limited HCN/CO linewidth ratio, accuracy of the FIR SED fit, AGN bias of $L_{\rm FIR}$). The statistical and individual results, so far, would even be consistent with an even stronger increase in $L_{\rm FIR}$/$L'_{\rm HCN}$ toward the highest $z$ and/or $L_{\rm FIR}$. Our study of J1148+5251 may hint at such an effect. Clearly, it is desirable to obtain more sensitive observations of this source to further investigate this issue. Due to its superior collecting area and high calibrational stability, the VLA is ideally suited for such a sensitive study. Although J1148+5251 is the most CO- and FIR-luminous $z$$>$6 galaxy known, 80\,hr of VLA observations were necessary to obtain the current limit. In a favourable case, the \bhcn\ line may have a strength of about 1.5 times the current rms. To obtain a solid 5\,$\sigma$ detection of such a line, of order 1000\,hr of observations with the VLA would be required. Due to improved receivers and antenna performance, the fully operational EVLA will be by a factor of two more sensitive to spectral lines of several 100\,\kms\ width (such as in J1148+5251), but will still require long integration times. Studies of dense gas at $z$$>$6 thus appear to require an order of magnitude increase in collecting area, such as offered by future facilities like the SKA phase I demonstrator (e.g., Carilli \citeyear{car06}), which can serve as a low frequency counterpart to the Atacama Large Millimeter/submillimeter Array (ALMA). | 7 | 10 | 0710.4525 |
0710 | 0710.4149_arXiv.txt | {The Solar Tower Atmospheric Cherenkov Effect Experiment (STACEE) is an atmospheric Cherenkov telescope (ACT) that uses a large mirror array to achieve a relatively low energy threshold. For sources with Crab-like spectra, at high elevations, the detector response peaks near 100 GeV. Gamma-ray burst (GRB) observations have been a high priority for the STACEE collaboration since the inception of the experiment. We present the results of 20 GRB follow-up observations at times ranging from 3 minutes to 15 hours after the burst triggers. Where redshift measurements are available, we place constraints on the intrinsic high-energy spectra of the bursts.} \begin{document} | The Solar Tower Atmospheric Cherenkov Effect Experiment (STACEE) is a showerfront-sampling Cherenkov telescope sensitive to gamma rays above 100 GeV. It is located at the National Solar Thermal Test Facility (NSTTF) at Sandia National Laboratories outside Albuquerque, New Mexico, USA. The NSTTF is located at 34.96$^{\circ}$N, 106.51$^{\circ}$W and is 1700 m above sea level. The facility has 220 heliostat mirrors designed to track the sun across the sky, each with 37 m$^{2}$ area. STACEE uses 64 of these heliostats to collect Cherenkov light produced by air showers. STACEE employs five secondary mirrors on the solar tower to focus the Cherenkov light onto photomultiplier tube (PMT) cameras, as shown in Figure \ref{concept}. The light from each heliostat is detected by a separate PMT and the waveform of the PMT signal is recorded by a flash ADC. A programmable digital delay and trigger system\cite{IEEENSS2000} selects showers for acquisition while eliminating most random coincidences of night sky background photons. Under good observing conditions, STACEE operates with a threshold around 5 photoelectrons per channel. A detailed description of the instrument can be found in D.M. Gingrich et al.\cite{Gingrich05}. \begin{figure}[t] \begin{minipage}{0.5\textwidth} \begin{center} \includegraphics[width=.9\textwidth]{concept5.eps} \end{center} \end{minipage} \hfill \begin{minipage}{0.45\textwidth} \caption{Conceptual drawing of STACEE.} \label{concept} \vskip 1in \end{minipage} \end{figure} The large mirror area of the STACEE detector leads to an energy threshold lower than those attainable by most single-dish imaging telescopes or water-Cherenkov telescopes. The energy threshold - defined as the energy at which the trigger rate peaks - is determined by the spectrum of the source and the effective area of the detector at the target position. For targets above 60$^{\circ}$ in elevation with power-law spectral indices between 2 and 3, the energy threshold is typically between 150 and 200 GeV. For targets near zenith, STACEE has significant effective area at energies as low as 50 GeV. A low energy threshold opens up the possibility of detecting more distant sources\cite{Primack05}. Collisions of gamma rays with extragalactic background light (EBL) photons produce electron-positron pairs, attenuating the gamma-ray flux from distant sources. The extinction becomes more severe with increasing energy, producing an energy-dependent horizon for gamma-ray observations. Thus, a low energy threshold is advantageous when attempting to characterize the high-energy emission of GRBs, for which the mean measured redshift (for Swift bursts) is 2.7\cite{Le07}. STACEE observations are typically taken in pairs of on-source and off-source runs. The off-source runs serve as measurements of the background event rate produced by cosmic-ray showers. Under normal conditions, the cosmic-ray trigger rate is $\sim$5 Hz. Background rejection techniques have been explored and the most effective technique is described elsewhere\cite{Kildea05,Kildea05_2,Lindner06}. After background rejection cuts, STACEE typically obtains a 5$\sigma$ detection of the Crab with approximately 10 hours of on-source observations. Under good observing conditions, STACEE would obtain a 5$\sigma$ detection in 30 minutes for a source with a spectrum equal to 4.5 times that of the Crab. | 7 | 10 | 0710.4149 |
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0710 | 0710.1841_arXiv.txt | We report the discovery of a planet transiting a moderately bright ($V=12.00$) G star, with an orbital period of $2.788491\pm0.000025$ days. From the transit light curve we determine that the radius of the planet is $\rpl = 1.257\pm0.053\,\rjuplong$. \hatcurb\ has a mass of $\mpl = \hatcurm\,\mjuplong$, similar to the average mass of previously-known transiting exoplanets, and a density of $\rhopl = \hatcurrho\,\gcmc$. We find that the center of transit is $T_{\mathrm{c}} = \hatcurT$ (HJD), and the total transit duration is \hatcurdur\,days. | \label{sec:intro} To date about 20 extrasolar planets have been found which transit their parent stars and thus yield values for their mass and radius\footnote{Extrasolar Planets Encyclopedia: http://exoplanet.eu}. Masses range from 0.07\,\mjup\ \citep[GJ436;][]{gillon07} to about 9\,\mjup\ \citep[HAT-P-2b;][]{bakos07}, and radii from 0.7\,\rjup\ (GJ436) to about 1.7\,\rjup\ \citep[TRES-4;][]{mandushev07}. These data provide an opportunity to compare observations with theoretical models of planetary structure across a wide range of parameters, including those of the host star \citep[e.g.][and references therein]{burrows07,fortney07}. Transits also yield precise determination of other physical parameters of the extrasolar planets, for instance the surface gravity. Interesting correlations between these parameters have been noted early on, such as that between masses and periods \citep{mazeh05} or periods and surface gravities \citep{southworth07}. Classes of these close-in planets have also been suggested, such as very hot Jupiters (VHJs; P=1--3 days) and hot Jupiters \citep[HJs; P=3--9 days;][]{gaudi05}, or a possible dichotomy based on Safronov numbers \citep{hansen07}. However, the small ensemble of transiting exoplanets (TEPs) does not allow robust conclusions, thus the addition of new discoveries is valuable. Over the past year the HATNet project\footnote{www.hatnet.hu} \citep{bakos02,bakos04}, a wide-angle photometric survey, has announced four TEPs. In this Letter we report on the detection of a new transiting exoplanet, which we label \hatcurb, and our determination of its parameters, such as mass, radius, density and surface gravity. | \label{sec:disc} \hatcurb\ is an ordinary hot Jupiter (P = 2.788 days) with slightly inflated radius (\rpl = 1.26\rjup) for its mass of 1.06\mjup, orbiting a slightly metal rich solar-like star. The $\sim$20\% radius inflation is what current models predict for a planet with equilibrium temperature of $\sim$1500K \citep{burrows07,fortney07}. \hatcurb\ is more massive than any of the known TEPs with similar period ($2.5\lesssim P \lesssim3$\,d), such as XO-2b, WASP-1, HAT-P-3b, TRES-1, and HAT-P-4b, with the exception of TRES-2. The latter is fairly similar in mass, radius, orbital period, and stellar effective temperature. However, \hatcurb\ is interesting in that it falls between Class I and II, as defined by the Safronov number and $T_{eq}$ of the planet \citep{hansen07}. \hatcurb\ has a Safronov number of $0.059\pm0.005$ , while Class I is defined as $0.070\pm0.01$, especially at $T_{eq}\sim 1500$K. It seems that the additional discovery and characterization of transiting planets of Jupiter and higher masses would be very helpful in order to understand these new correlations and their reality. | 7 | 10 | 0710.1841 |
0710 | 0710.1246_arXiv.txt | We present an algorithm for solving the linear dispersion relation in an inhomogeneous, magnetised, relativistic plasma. The method is a generalisation of a previously reported algorithm that was limited to the homogeneous case. The extension involves projecting the spatial dependence of the perturbations onto a set of basis functions that satisfy the boundary conditions (spectral Galerkin method). To test this algorithm in the homogeneous case, we derive an analytical expression for the growth rate of the Weibel instability for a relativistic Maxwellian distribution and compare it with the numerical results. In the inhomogeneous case, we present solutions of the dispersion relation for the relativistic tearing mode, making no assumption about the thickness of the current sheet, and check the numerical method against the analytical expression. | Dissipation of the energy carried by relativistic plasma outflows is important for the physics of pulsar winds and gamma-ray bursts (for recent reviews see \cite{2007astro.ph..3116K} and \cite{2005RvMP...76.1143P}). In these objects, the plasma is probably composed of electrons, positrons and protons. As well as being in relativistic bulk motion with respect to the observer, the random thermal energy of the plasma may also be relativistic, i.e., comparable to the rest mass energy of the constituent particles. In this paper we concentrate on two instabilities: these are the two-stream or Weibel instability~\cite{1959PhRvL...2...83W} and the tearing modes in a relativistic neutral pair plasma current sheet, thought to play a role in the formation process of relativistic shocks \cite{2007arXiv0706.3126S} and the dissipation of magnetic energy in pulsar winds \cite{2005PPCF...47B.719K}. They are investigated by generalising and extending to the inhomogeneous case the method presented in \cite{2007PPCF...49..297P}. Motivated primarily by the need for code verification, we have derived some analytical expressions for the linear growth rates of these instabilities. The Weibel instability is very important in astrophysical processes because it is able to generate a magnetic field by extracting the free energy from an anisotropic momentum distribution in an unmagnetised plasma or from the kinetic drift energy. There is an extensive literature on the Weibel instability: general conditions for the existence of the relativistic Weibel instability for arbitrary distribution functions are discussed in~\cite{2006PhPl...13b2107S}, and wave propagation in counter-streaming magnetised nonrelativistic Maxwellian plasmas are studied in~\cite{2005PhPl...12.2901T,2006PhPl...13f2901T}. Dispersion curves have been found in some special cases such as, for example, the fully relativistic bi-Maxwellian distribution function, (Yoon~\cite{1989PhFlB...1.1336Y}), and the water-bag distribution function, in which case closed-form analytical expressions can be derived not only for the Weibel instability~(Yoon~\cite{1987PhRvA..35.2718Y}), but also for the cyclotron maser and whistler instabilities (Yoon~\cite{1987PhRvA..35.2619Y}). However, finding an analytical expression for the dispersion relation for a given equilibrium distribution function is a complicated or even impossible task. It involves a four-dimensional integration (3D in momentum space and 1D in time) of the equilibrium distribution function which is difficult to perform in closed form. For this reason, the water-bag distribution is the preferred profile to analyse magnetic field generation in fast ignitor scenarios, (Silva et al.~\cite{2002PhPl....9.2458S}) and in relativistic shocks, (Wiersma and Achterberg~\cite{2004A&A...428..365W}, Lyubarsky and Eichler~\cite{2006ApJ...647.1250L}). The Weibel instability in a magnetised electron-positron pair plasma has been investigated by Yang et al~\cite{1993PhFlB...5.3369Y} using two model distributions: the water bag, and one with a power-law dependence at high energy. A general covariant description has been formulated by Melrose~\cite{1982AuJPh..35...41M} and by Schlickeiser~\cite{2004PhPl...11g5532S}. In the present work, we focus on equilibrium configurations given by a relativistic Maxwellian distribution function, which allows one to reduce the four-dimensional integral to a simple one-dimensional integral, as we demonstrate in Section~\ref{sec:weibel_analytical}. The growth rates are then found by solving this equation using a single numerical quadrature, and are compared to the results found using our extended algorithm in Section~\ref{sec:weibel_numerical}. In the inhomogeneous case, the stability properties of a nonrelativistic Harris current sheet also have a substantial literature, with notable recent studies by Daughton~\cite{1999PhPl....6.1329D} and Silin et al.~\cite{2002PhPl....9.1104S}. In the relativistic case, the tearing mode instability has been investigated by Zelenyi \& Krasnoselskikh~\cite{1979SvA....23..460Z}, by integrating first order perturbations of the relativistic Maxwellian distribution function along approximate, straight-line particle trajectories, in the thick layer limit (in which the Larmor radius is much smaller than the thickness of the current sheet). In Section~\ref{sec:tearing_analytical} we lift these restrictions to present new results for the tearing mode instability in a neutral current sheet of arbitrary temperature and thickness, and compare these with the results found using the generalised algorithm in Section~\ref{sec:tearing_numerical}. In this work, no assumption is made about the thickness of the current sheet, and the particle trajectories are found numerically in the background magnetic field. The numerical method, which is an extension of our previous algorithm, \cite{2007PPCF...49..297P}, that computes the linear dispersion relation of waves within a Vlasov-Maxwell description, is described in Section~\ref{sec:algorithm}. It is based on the approach of Daughton~\cite{1999PhPl....6.1329D} for non relativistic Maxwellians, and involves explicit time integration of particle orbits along the unperturbed trajectories. We modify and extend our former code to include inhomogeneities in the plasma equilibrium configuration. Moreover, we generalise it to a fully relativistic approach, i.e., one that allows for relativistic temperatures as well as relativistic drift speeds. | \label{sec:conclusion} We present a generalisation and extension of our previous algorithm~\cite{2007PPCF...49..297P} to solve the linear dispersion relation for relativistic multi-component inhomogeneous and magnetised plasmas. The code is validated by comparing the results with two standard configurations: the relativistic Weibel instability in a homogeneous plasma, and the tearing mode instability in a relativistic neutral Harris sheet. To effect the comparison, we derived useful analytical expressions, Eq.~(\ref{eq:RelDispWeibel}) and Eq.~(\ref{eq:Harris2}), for the dispersion relations in these configurations and solved them numerically. We conclude that this code is a suitable tool for the study of stability properties of more general configurations of interest in gamma-ray burst and pulsar wind theories. \ack{This research was supported by a grant from the G.I.F., the German-Israeli Foundation for Scientific Research and Development.} | 7 | 10 | 0710.1246 |
0710 | 0710.3769_arXiv.txt | {We describe a broad class of time-dependent exact wave solutions to 6D gauged chiral supergravity with two compact dimensions. These 6D solutions are nontrivial warped generalizations of 4D pp-waves and Kundt class solutions and describe how a broad class of previously-static compactifications from 6D to 4D (sourced by two 3-branes) respond to waves moving along one of the uncompactified directions. Because our methods are generally applicable to any higher dimensional supergravity they are likely to be of use for finding the supergravity limit of time-dependent solutions in string theory. The 6D solutions are interesting in their own right, describing 6D shock waves induced by high energy particles on the branes, and as descriptions of the near-brane limit of the transient wavefront arising from a local bubble-nucleation event on one of the branes, such as might occur if a tension-changing phase transition were to occur.} \preprint{PI-COSMO-65} \begin{document} | Understanding time-dependent dynamics is central to applications of higher-dimensional supergravity theories \cite{HiDSugra} to cosmology and to particle physics. Inasmuch as higher-dimensional supergravities provide the low-energy limit of string theories, any understanding of time-dependence in the supergravity limit also provides a guide for the thornier issue of understanding these same issues in string theory. For these reasons there is considerable interest in finding time-dependent solutions to higher-dimensional supergravity \cite{tdepsusy} (as well as of non-supersymmetric gravity \cite{tdepnonsusy}, since this can also sometimes capture similar physics). Much of this activity has focussed on 10D and 5D theories (motivated by string applications, and Randall-Sundrum \cite{RS} constructions), although these can either be more difficult to solve or they can have features which are specific to the relative simplicity of co-dimension one spaces. Six-dimensional supergravity has more recently emerged as being a useful intermediate workshop within which to investigate phenomena which can often generalize to still higher-dimensional contexts. Interest in 6D supergravity has been further sharpened by the recognition that it can provide insights into the nature of the cosmological constant problem \cite{SLED1}--\cite{SLEDpheno}, by building on the observation that higher-dimensional theories can break the link between the 4D vacuum energy density and the curvature of 4D spacetime \cite{5DSelfTune}--\cite{6DNonSUSYSelfTunex} (see \cite{Burgess:2007ui} for a review). There has also been considerable recent interest in 6 dimensional models more generally \cite{Closelyrelatedworks}. Including branes is notoriously difficult due the necessity to regularize UV divergences which arise \cite{regularizations}, however for recent work on understanding this more deeply in the context of effective field theory see \cite{claudiaeft}. In this paper we further the program of understanding time-dependent solutions to higher-dimensional supergravity in two ways. In \S2, we present a method for constructing explicit exact solutions to the supergravity field equations of 6D gauged chiral supergravity, to derive a new class of exact solutions to these equations which describe a class of gravitational waves passing through spacetimes for which two dimensions are compactified in response to the stress-energy of two space-filling 3-branes. In \S3 we discuss the applications of these solutions, including how to construct the shock wave metric corresponding to ultra-relativistic particles and BH's moving on one of the branes, and also how these solutions provide some insight into the transient part of the dynamics describing outgoing waves which would arise shortly after a phase transition on one of the source branes. | Six dimensional supergravity provides a fruitful laboratory for investigating the issues which underly higher-dimensional physics in general, and brane approaches to the cosmological constant problem in particular. It does so because 6D is rich enough to exhibit many of the properties of still-higher dimensions --- like moduli-stabilization through fluxes \cite{Susha}, brane back-reaction on internal geometries, chiral fermions and Green-Schwarz anomaly cancellation \cite{6Danomalies}. Yet it is also on the one hand simple enough to allow the development of techniques for obtaining physically interesting exact solutions, but on the other hand not so simple as to be misleading about what happens in higher dimensional (in a way which co-dimension one physics sometimes can be). We have used these properties to explore solution-generation techniques which we believe to be applicable to a wide variety of higher-dimensional supergravities. We do so by using these techniques to construct a new class of time-dependent exact solutions to the field equations of 6D chiral gauged supergravity. These solutions describe the physics of nonlinear gravitational waves passing through compactified spacetimes for which two dimensions are self-consistently compactified in response to the presence of two space-filling branes and a bulk Maxwell flux. We believe these methods to merit more detailed exploration. | 7 | 10 | 0710.3769 |
0710 | 0710.0901_arXiv.txt | The diffuse interstellar bands (DIBs) probably arise from complex organic molecules whose strength in local galaxies correlates with neutral hydrogen column density, \nhi, and dust reddening, \ebmv. Since Ca{\sc \,ii} absorbers in quasar (QSO) spectra are posited to have high \nhi\ and significant \ebmv, they represent promising sites for the detection of DIBs at cosmological distances. Here we present the results from the first search for DIBs in 9 Ca{\sc \,ii}-selected absorbers at $0.07 < z_{\rm abs} < 0.55$. We detect the 5780\,\AA\ DIB in one line of sight at $z_{\rm abs} = 0.1556$; this is only the second QSO absorber in which a DIB has been detected. Unlike the majority of local DIB sight-lines, both QSO absorbers with detected DIBs show weak 6284\,\AA\ absorption compared with the 5780\,\AA\ band. This may be indicative of different physical conditions in intermediate redshift QSO absorbers compared with local galaxies. Assuming that local relations between the 5780\,\AA\ DIB strength and \nhi\ and \ebmv\ apply in QSO absorbers, DIB detections and limits can be used to derive \nhi\ and \ebmv. For the one absorber in this study with a detected DIB, we derive \ebmv\ = 0.23\,mag and $\log$\nhi\ $\ge$ 20.9, consistent with previous conclusions that Ca{\sc \,ii} systems have high H{\sc \,i} column densities and significant reddening. For the remaining 8 Ca{\sc \,ii}-selected absorbers with 5780\,\AA\ DIB non-detections, we derive \ebmv\ upper limits of 0.1--0.3\,mag. | Damped Lyman alpha (DLA) systems are usually considered to be the class of QSO absorber with the highest neutral hydrogen column densities [\nhi\ $\ge 2 \times 10^{20}$ \cm]. Nonetheless, the DLAs are characterised by generally low metallicities and gas-phase depletion fractions \citep[e.g.][]{KhareP_04a,AkermanC_05a,ProchaskaJ_07a} and low reddening due to dust (\citealt{MurphyM_04c}; \citealt*{EllisonS_05a}). The DLAs are also poor in molecules, as demonstrated by both their generally low fractions of H$_2$ \citep[e.g.][]{LedouxC_03a} and the lack of a detection for any other molecular species, such as OH or CO \citep[e.g.][]{CurranS_06a}. Although the handful of DLAs which do exhibit molecular H$_2$ absorption may be biased, e.g. towards high metallicities \citep{PetitjeanP_06a}, such systems can offer a novel insight into the physical conditions of the galactic interstellar medium \citep[ISM, e.g.][]{SrianandR_05a,NoterdaemeP_07a}. In addition to the study of H$_2$, one avenue that is just starting to be explored is how the diffuse interstellar bands \citep[DIBs; see reviews by][]{HerbigG_95a,SarreP_06a} may be used to probe the intermediate redshift ISM. Although lacking definitive identifications, the strength (both absolute and relative) of these broad absorption features in the Milky Way (MW) and other nearby galaxies exhibit dependencies on (and sometimes, tight correlations with) neutral gas content, dust reddening, metallicity and local radiation field \citep[e.g.][]{HerbigG_93a,CoxN_06a,WeltyD_06a,CoxN_07a}. Moreover, if DIBs are as strong in DLAs as they are in the MW [i.e.~for a given \nhi], then they should be relatively easy to detect at intermediate redshifts. The first systematic search for DIBs in DLAs has recently been carried out by Lawton et al.~(in preparation) in 7 $z < 1$ absorbers. In only one case were DIBs detected: the 4428, 5705 and 5780\,\AA\ features\footnote{We cite all DIBs with reference to their normal air wavelengths, although their vacuum values have been used in practice in order to be consistent with our spectral wavelength calibration; see Section 2.} were all detected in the $z \sim 0.5$ DLA towards AO 0235$+$164 \citep{JunkkarinenV_04a,YorkB_06a}. Lawton et al.~showed that for the 6 non-detections in their DLA sample, the strength of the 5780\,\AA\ DIB [which shows one of the tightest correlations with \nhi\ in the MW] is often at least 3 times weaker in DLAs for a given \nhi\ compared with Galactic sight-lines. The 6284\,\AA\ DIB is even more under-abundant in DLAs for a given \nhi: 4--10 times weaker than towards Galactic sight-lines. A similar result has been found for DIBs in the Large and Small Magellanic Cloud \citep[LMC and SMC;][]{WeltyD_06a} where the 5780\,\AA\ DIB is typically 10--30 times weaker than expected from the Galactic relation. On the other hand, the 5780\,\AA\ DIB strength correlates well with \ebmv\ in both Galactic and Magellanic Cloud sight-lines, and the detection towards AO 0235$+$164 also lies on the same relationship \citep{YorkB_06a}. These results hint that DIB formation/survival and high dust content are closely linked and that DIBs are therefore most likely to be detected in galaxies with high reddening. \citet*{WildV_06a} have recently suggested that absorbers identified via high equivalent widths (EWs) of Ca{\sc \,ii} may select the highest \nhi\ and highest \ebmv\ absorbers. For example, whereas DLAs have been constrained to have \ebmv\ $\le$ 0.04 \citep{MurphyM_04c,EllisonS_05a}, \citet{WildV_06a} find that absorbers with Ca{\sc \,ii} $\lambda$3934 EWs $>$0.7\,\AA\ have \ebmv\ values up to $\sim$0.1 mag. Ca{\sc \,ii} absorbers may therefore be promising sites for the detection of DIBs. | \begin{figure} \centerline{\rotatebox{0}{\resizebox{9cm}{!} {\includegraphics{EBMV_NaI_5780.ps}}}} \caption{\label{ebmv_fig} Correlations of Na{\sc \,i} (top panel), \nhi\ (middle panel) and \ebmv\ (bottom panel) versus the log of the EW (in m\AA) of the 5780\,\AA\ DIB. Open squares are Galactic points from \citet{HerbigG_93a}; open triangles/diamonds are SMC/LMC points from \citet{WeltyD_06a}, \citet{VladiloG_87a} and \citet{Vidal-MadjarA_87a}; crosses are other extra-galactic data points taken from \citet{SollermanJ_05a}, \citet{DOdoricoS_89a} and \citet{HeckmanT_00b}; solid stars are DLAs from \citet{YorkB_06a} and Lawton et al.~(in preparation). The best (least squares) fit to the data in each panel is shown with a solid line and the fit solution given at the bottom of each panel. The fit of the 5780\,\AA\ DIB with \ebmv\ uses all available data; the fits with of 5780 with \nhi\ and Na{\sc \,i} use just the Galactic data. In the middle and lower panels, the solid vertical line indicates the 5780\,\AA\ DIB detection towards J0013$-$0024 and the dotted lines are the 3$\sigma$ upper limits for 6 other sight-lines where we have a 5780\,\AA\ limits (as given in Table \ref{dib_table}). } \end{figure} In local (e.g.~Galactic, LMC, SMC) sight-lines, the 6284\,\AA\ DIB is typically 2--3 times stronger than the 5780\,\AA\ DIB (e.g., York et al. 2006a and references therein). The one exception is the unusual SMC wing sight-line towards Sk~143 where the 6284\,\AA\ DIB has an EW less than half that of the 5780\,\AA\ DIB \citep{WeltyD_06a}. \citet{YorkB_06a} also found that in the one DLA sight-line with a 5780\,\AA\ band detection out of the 7 studied by Lawton et al.~(in preparation), the EW of the 6284\,\AA\ feature was also constrained to be less than the EW of the 5780\,\AA\ line. \citet{YorkB_06a} suggested that these unusual line ratios could be an indication of an ISM that is more protected from the ambient UV radiation field. In the Ca{\sc \,ii}-selected absorber towards J0013$-$0024, we constrain the EW of the 6284\,\AA\ DIB to be at least $\sim$20\% weaker than the 5780\,\AA\ band. The DIB ratios in this absorber are therefore consistent with those in the DLA detection of \citet{YorkB_06a} and the SMC wing sight-line Sk~143 but inconsistent with other local sight-lines, including starburst galaxies \citep{HeckmanT_00b} and the Magellanic Clouds \citep{WeltyD_06a}. In the Galaxy, many DIBs show correlations with \nhi\ and $N$(Na{\sc \,i}) \citep[e.g.][]{HerbigG_93a,HerbigG_95a}. In Table \ref{dib_table} we tabulate the EWs of the Na{\sc \,i} doublet for our Ca{\sc \,ii}-selected absorbers. However, we do not calculate the Na{\sc \,i} column density because, if the lines are strong enough to be detected in our moderate resolution spectra, they are likely to be saturated. The Galactic correlations of \nhi\ and $N$(Na{\sc \,i}) with the 5780\,\AA\ DIB, which is one of the tightest of the DIB relations, is shown in Fig. \ref{ebmv_fig}. We also show data for the Magellanic Clouds \citep{WeltyD_06a} and DLAs (\citealt{YorkB_06a}; Lawton et al.~in preparation), where it can be seen that the DIBs are weak for their \nhi\ compared with the Galactic correlation. As shown in Fig. \ref{ebmv_fig}, the DIBs in extra-galactic sight-lines are also weak for their Na{\sc \,i} column densities. These departures from the Galactic relations are probably due to a combination of effects including ambient radiation field, metallicity and dust-to-gas ratios \citep{CoxN_06a}. Assuming that the Galactic 5780--\nhi\ relation provides a lower limit for the H{\sc \,i} column density, DIB detections may be useful for constraining \nhi\ in the absence of \lya\ observations. For example, \citet{WildV_05a} and \citet{WildV_06a} have argued that Ca{\sc \,ii} absorbers represent the high column density end of the DLA distribution. Our detection of the 5780\,\AA\ DIB in the $z_{\rm abs} = 0.1556$ absorber towards J0013$-$0024 supports this hypothesis, and we derive $\log$\nhi\ $\ge$ 20.9 for this absorber. Unlike correlations with \nhi\ and $N$(Na{\sc \,i}), \citet{WeltyD_06a} have shown that the 5780\,\AA\ DIB strength follows a single relationship with \ebmv\ in both Galactic and Magellanic Cloud sight-lines. \citet{YorkB_06a} found that the single DLA 5780\,\AA\ DIB detection towards AO 0235$+$164 fell on the same relationship. It is not yet clear whether the apparent universality of this correlation is driven by a tight physical connection between dust properties and DIB formation \citep{CoxN_07a} or whether it is coincidence of different physical drivers working in different directions \citep{CoxN_06a}. However, if the 5780--\ebmv\ is applicable to QSO absorbers, we can use our DIB detection limits to constrain their reddening. \citet{WeltyD_06a} derive a best fit correlation between the 5780\,\AA\ DIB (in m\AA) and the \ebmv\ for Galactic sight-lines: log \ebmv\ = $-$2.70 + 1.01 log EW(5780). We derive the best fit relation to the 5780--\ebmv\ data points of the Galactic plus Magellanic Cloud plus AO 0235$+$164 DLA sight-lines and find log \ebmv\ = $-$2.19 + 0.79 log EW(5780) (see Figure \ref{ebmv_fig}). The range in log \ebmv\ values around the best fit relation is $\sim \pm$ 0.4 dex. This correlation gives a reddening for the Ca{\sc \,ii} absorber towards J0013$-$0024 of \ebmv$\,\,\sim0.23$\,mag and upper limits for the other 8 Ca{\sc \,ii} absorbers in our sample of 0.1--0.3\,mag. These values provide independent estimates of reddening associated with Ca{\sc \,ii}-selected absorbers that do not depend directly on the choice of extinction law and can be applied for individual absorbers and not just in a statistical fashion \citep[e.g.][]{MurphyM_04c,WildV_05a,WildV_06a}. The Ca{\sc \,ii} EWs of our sample are typically $<0.7$\,\AA\ (see Table \ref{dib_table}); for this range of EWs, \citet{WildV_06a} determine average reddenings of \ebmv\ = 0.02, 0.03 and 0.03\,mag for MW, LMC and SMC extinction curves respectively. | 7 | 10 | 0710.0901 |
0710 | 0710.5523_arXiv.txt | We have performed a high mass and force resolution simulation of an idealized galaxy forming from dissipational collapse of gas embedded in a spherical dark matter halo. The simulation includes star formation and effects of stellar feedback. In our simulation a stellar disk forms with a surface density profile consisting of an inner exponential breaking to a steeper outer exponential. The break forms early on and persists throughout the evolution, moving outward as more gas is able to cool and add mass to the disk. The parameters of the break are in excellent agreement with observations. The break corresponds with a rapid drop in the star formation rate associated with a drop in the cooled gas surface density, but the outer exponential is populated by stars that were scattered outward on nearly circular orbits from the inner disk by spiral arms. The resulting profile and its associated break are therefore a consequence of the interplay between a radial star formation cutoff and redistribution of stellar mass by secular processes. A consequence of such evolution is a sharp change in the radial mean stellar age profile at the break radius. | \label{sec:intro} Since the early work of \citet{de-Vaucouleurs:1958} it has been recognized that the disks of spiral galaxies generally follow an exponential radial surface brightness profile, and various theories have explored the physical causes and consequences of this property \citep[e.g.][]{Fall:1980lr, Lin:1987pb, Dalcanton:1997bh, Mo:1998mi, van-den-Bosch:2001aa}. However, since \citet{van-der-Kruit:1979gb, van-der-Kruit:1987fk} it has been known that the outer regions of disks exhibit more varied behavior. This has been confirmed by an abundance of recent data \citep[e.g.,][hereafter PT06]{Pohlen:2000ff, Pohlen:2002ec, Bland-Hawthorn:2005ys, Erwin:2005kl, Pohlen:2006lh}. In a sample of nearby late-type galaxies from the Sloan Digital Sky Survey, PT06 found that about 60\% have an inner exponential followed by a steeper outer exponential (downward-bending), $\sim$30\% have the inner exponential followed by a shallower outer exponential (upward-bending), while only $\sim10$\% have no detectable breaks. Therefore breaks are a common feature of disk galaxies that any complete theory of galaxy formation must be able to explain. Furthermore, the discovery of UV emission at radii well beyond the H$\alpha$ emission cutoff \citep{Gil-de-Paz:2005, Thilker:2005, Thilker:2007a}, the observational evidence for inside-out disk growth \citep{Munoz-Mateos:2007}, and detections of disk breaks in the distant universe \citep{Perez:2004, Trujillo:2005}, suggest that the outer disks provide a direct view of disk assembly. Several theories for the formation of breaks have been investigated. \Citet{van-der-Kruit:1987fk} proposed that angular momentum conservation in a collapsing, uniformly rotating cloud naturally gives rise to disk breaks at roughly 4.5 scale radii. \Citet{van-den-Bosch:2001aa} suggested that breaks are due to angular momentum cutoffs of the cooled gas. More commonly breaks have been attributed to a threshold for star formation (SF), whether due to low gas density which stabilizes the disk \citep{Kennicutt:1989bs}, or to a lack of a cool equilibrium ISM phase \citep{Elmegreen:1994eu, Schaye:2004kb}. Using a semi-analytic model, \citet{Elmegreen:2006oq} demonstrate that a double-exponential profile may result from a multi-component star formation prescription. The existence of extended UV disks \citep[e.g][]{Thilker:2007a} and the lack of a clear correlation of H$\alpha$ cut-offs and optical disk breaks \citep{Pohlen:2004, Hunter:2006aa} further complicates the picture. Regardless, while a sharp SF cutoff may explain the disk truncation, it does not provide a compelling explanation for extended outer exponential components. Alternatively, \citet{Debattista:2006wd} demonstrated that the redistribution of angular momentum by spirals during bar formation also produces realistic breaks in collisionless $N$-body simulations. In this Letter we present the first results from a series of high-resolution Smooth Particle Hydrodynamics (SPH) simulations of isolated galaxy formation aimed at exploring the formation and evolution of disk breaks and outskirts in a massive, high surface brightness galaxy without a strong central bar. Resulting breaks are analogous to downward-bending breaks seen in observations. The clear advantage of our approach over past attempts is that we use a fully self-consistent physical model of the system, making no a priori assumptions about the distribution of material in the disk. Rather, we allow the disk to grow spontaneously under the effects of gravity and gas hydrodynamics, itself influenced by star formation and feedback. The $N$-body approach (at sufficiently high mass and force resolution) ensures that we capture the dynamical processes contributing to disk evolution. Furthermore, the inclusion of prescriptions for SF and feedback allows us to make observational predictions across the break region. We show that (1) the break forms rapidly ($\lesssim$ 1Gyr) and persists throughout the evolution of the system, moving outward as the disk mass grows; (2) the break is seeded by a sharp decrease in star formation which is caused in our simulation by a rapid decrease in the surface density of cool gas; (3) the outer disk is populated by stars that have migrated, on nearly circular orbits, from the inner disk, and consequently the break is associated with a sharp change in the radial mean stellar age profile; (4) break parameters agree with current observations. | \label{sec:conclusions} We have shown that in a self-consistent model, where the stellar disk forms through gas cooling and subsequent star formation within a dark matter halo, breaks in the stellar surface density form through the combination of different effects. A rapid drop in the SFR seeds the break and secular evolution populates the outer exponential. In our model the SFR drop is due wholly to a drop in the surface density of gas. However a break in SFR induced by other means ({\it e.g.} a volume density threshold or perhaps warps) would lead to similar behavior of the outer disk and stellar density break parameters. Our model properties satisfy current observational constraints, both in the statistical sense of break properties from galaxies in SDSS, as well as the much more detailed observations of breaks in stellar populations of NGC~4244. Though our model does not account for the effects of evolution in a full cosmological context, its simplicity assures that this is the minimal degree of evolution (with no interactions or formation of a significant bar) and should therefore be rather generic provided that a disk is massive enough for strong transient spirals to form. The model predicts that there should be an abrupt change in the radial mean stellar age profile coincident with the break, which can be tested with future observations. | 7 | 10 | 0710.5523 |
0710 | 0710.2135_arXiv.txt | IC 4406 is a large (about 100$''$ $\times$ 30$''$) southern bipolar planetary nebula, composed of two elongated lobes extending from a bright central region, where there is evidence for the presence of a large torus of gas and dust. We show new observations of this source performed with IRAC (Spitzer Space Telescope) and the Australia Telescope Compact Array. The radio maps show that the flux from the ionized gas is concentrated in the bright central region and originates in a clumpy structure previously observed in H$\alpha$, while in the infrared images filaments and clumps can be seen in the extended nebular envelope, the central region showing toroidal emission. Modeling of the infrared emission leads to the conclusion that several dust components are present in the nebula. | \label{introduction} IC 4406 is a well-studied southern planetary nebula. It has been imaged with several telescopes at different wavelength ranges. Near-IR images show two H$_2$ lobes \citep{storey}, orthogonal to the nebula's major axis and $\sim$25$''$ away from each other. These peaks are approximately coincident with the two blobs observed in H$\alpha$+[\ion{N}{2}] and [\ion{O}{3}] \citep{sahai}, interpreted as indicative of the presence of a dense equatorial torus of dust. The optical images show a central ionized region about 32$''$ in diameter. CO maps show the presence of a collimated high velocity outflow in the polar direction and with [CO]/[H$_2$]$\approx 5\times 10^{-6}$ and a total molecular mass in the range 0.16--3.2 M$_\odot$ \citep{sahai}. Hubble Space Telescope (HST) WFPC2 images in [\ion{N}{2}], H$\alpha$ and [\ion{O}{3}] have revealed the existence of an intricate system of dark lane features, which led to the name of \lq\lq Retina Nebula\rq\rq~for this object \citep{odell}. The nebula appears to be chemically homogeneous, as \citet{corradi} found no evidence of radial variation for He, O, N, Ne, and Ar. \citet{cox} have detected several C-rich features at mm wavelengths, such as CN, HCO$^+$, HCN and HNC, which indicate the nebula is C-rich, although a C/O ratio of 0.6 is reported by \citet{cohen}. IC 4406 is a relatively low electron density nebula. Values in the 400-2000 cm$^{-3}$ range have been estimated using several different optical and infrared lines, with values derived by [\ion{S}{2}] and [\ion{O}{3}] doublets matching around 540 cm$^{-3}$ \citep{liu,wang}. Its central star has a \ion{He}{2} Zanstra temperature of 96800 K \citep{phillips} and its distance is probably around 1.6 kpc \citep{sahai}, although some authors claim it may be overestimated \citep{odell}. \citet{gruenwald} have modeled IC 4406 with a 3-D photoionization code and fit many observed line intensities assuming there is a torus around the central star. They find as a best fit a central star temperature of $8\times10^4$ K, luminosity of 400 L$_\odot$, torus density 1500 cm$^{-3}$ and nebular density 100 cm$^{-3}$. In general, comparisons of IR images of planetary nebulae, which trace the molecular gas and warm dust emission, to optical line images, which trace the ionized gas, have shown the presence of similar structures \citep{latter}, leading to the conclusion that molecular and ionized gas spatially coexist in planetary nebulae, as well as dust grains, despite the different physical conditions these components are presumed to survive in. We have observed IC 4406 in the radio range to inspect the distribution of the ionized gas in its envelope and in the infrared to check for emission from the equatorial dust and molecular gas. In \S\ref{observations} we explain how we performed our observations and reduced the data; in \S\ref{results} we show our results and in particular in \S\ref{sed} how we modeled the emission in the radio and infrared ranges; \S\ref{nebular} compares our model results to the nebular parameter values obtained directly from the observational data; in \S\ref{summary} we summarize the present work. | \label{summary} We have observed IC 4406 in the cm and 3--10 micron ranges. Our radio observations have confirmed the presence of the complicated maze of lanes already observed in H$\alpha$ in the central region of the nebula and have enabled us to calculate several nebular parameters, whose values match the classification for this target as an evolved planetary nebula, in particular its low dust to gas mass ratio and density. IRAC imaging has revealed the presence of filaments in the nebula that were not detected in previous observations. Our IRAC measurements, combined with literature data at longer and shorter wavelengths, have enabled us to study the SED of the PN IC 4406 and reproduce it with DUSTY. This has revealed that three different dust components are needed to model the data, with temperatures ranging from 57 to 700 K. It has also been necessary to include in the model slightly larger grains than in the standard MRN composition (up to 6.5 $\mu$m) to account for the calculated 60 $\mu$m optical depth. The main limits of our modeled curve are the spherical geometry assumed in DUSTY and the lack of data in the mm and sub-mm ranges, which would give a constraint on the slope of the curve. As we have observed during our trials with DUSTY, the slope of the SED in the sub-mm range changes with the maximum size of the grains included in the model. Unfortunately, in this range observations are available only for a few stars so far: none for our target. We can speculate that in such a diversified dust environment, as we find in IC 4406, further lower temperature components may exist and future high sensitivity, high angular resolution observations will give a fundamental contribution to understand the physics of circumstellar envelopes in planetary nebulae. | 7 | 10 | 0710.2135 |
0710 | 0710.2904_arXiv.txt | Compared to planets around Sun--like stars, relatively little is known about the occurrence rate and orbital properties of planets around stars more massive than 1.3~\msun. The apparent deficit of planets around massive stars is due to a strong selection bias against early--type dwarfs in Doppler--based planet searches. One method to circumvent the difficulties inherent to massive main--sequence stars is to instead observe them after they have evolved onto the subgiant branch. We show how the cooler atmospheres and slower rotation velocities of subgiants make them ideal proxies for F-- and A--type stars. We present the early results from our planet search that reveal a paucity of planets orbiting within 1~AU of stars more massive than 1.5~\msun, and evidence of a rising trend in giant planet occurrence with stellar mass. | A planet--bearing star can be thought of as a very bright, extremely dense remnant of a protoplanetary disk. After all, a star inherits its defining characteristics---its mass and chemical composition---from the same disk material that forms its planets. The physical characteristics of planet host stars therefore provide a crucial link between the planets we detect today and the circumstellar environments from which they formed long ago. Studying the relationships between the observed occurrence rate and orbital properties of planets as a function of stellar characteristics informs theories of planet formation, and also helps guide the target selection of future planet searches. A wealth of recent work has demonstrated that planet occurrence is strongly correlated with chemical composition \citep{gonzalez97, santos04}; metal--rich stars are 3 times more likely to host planetary companions compared to stars with solar abundances \citep{fischer05b}. This finding can be understood in the context of the core accretion model. Increasing the metallicity of a star/disk system increases the surface density of solid material at the disk midplane, which in turn leads to an enhanced growth rate for protoplanetary cores \citep{ida04b, kornet05}. \begin{figure}[ht!] \plotone{mass_hist.eps} \caption {\footnotesize{Distribution of stellar masses for the target stars of the California and Carnegie Planet Search. The majority of the stars have masses between 0.7~\msun\ and 1.3~\msun. \label{mass_hist}}} \end{figure} Another factor that enhances the surface density of material in the disk midplane is its total mass. If the mass of circumstellar disks scales with the mass of the central star, then there should be an observed correlation between planet occurrence and stellar mass \citep{laughlin04, ida05b, kennedy07}. In principle, testing this hypothesis is fairly simple: one need only measure the fraction of stars with planets over a wide range of stellar masses. However, in practice such a study is not so straight forward given the limited range of stellar masses encompassed by most planet searches. The difficulty can be seen in Figure~\ref{mass_hist}, which shows the distribution of stellar masses for the target stars in California and Carnegie Planet Search \citep[CCPS; ][]{valenti05}, which is representative of most Doppler--based planet searches. Most of the stars in Figure~\ref{mass_hist} have masses between 0.7~\msun\ and 1.3~\msun. In a decidedly non--Copernican twist of nature, it turns out that stars like our Sun are ideal planet search targets. Solar--mass G and K dwarfs are slow rotators, have stable atmospheres, and are relatively bright. The fall--off toward lower stellar masses is simply because late K-- and M--type dwarfs are faint, making the acquisition of high--resolution spectra difficult without large telescope apertures \citep{butler06b, bonfils05b, endl03}. The sharp drop at higher stellar masses is due to a separate observational bias. Stars with spectral types earlier than F8 tend to have rotationally broadened spectral features ($V\sin{i} > 50$~\ks; Galland et al. 2005), have fewer spectral lines due to high surface temperatures, and display a large amount of atmospheric ``jitter.'' Stellar jitter is excess velocity scatter due to surface inhomogeneities and pulsation, which can approach 50--100~\ms\ for mid--F stars \citep{saar98, wright05}. These features conspire to limit the attainable radial velocity precision for early--type dwarfs to $> 50$~\ms, rendering exceedingly difficult the detection of all but the most massive and short--period planets. One method to circumvent the observational limitations inherent to high--mass dwarfs is to observe these stars after they have evolved away from the main--sequence. After stars have expended their core hydrogen fuel sources their radii expand, and their atmospheres cool leading to an increase in the number of metal lines in the star's spectrum. Stars crossing the subgiant branch also shed a large amount of angular momentum through the coupling of stellar winds to rotationally generated magnetic fields \citep{gray85, schrijver93, donascimento00}. The cooler atmospheres and slower rotational velocities of evolved stars lead to an increased number of narrow absorption lines in their spectra, making them much better suited for precision Doppler surveys. | 7 | 10 | 0710.2904 |
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0710 | 0710.4133_arXiv.txt | If light scalar fields are present at the end of inflation, their non-equilibrium dynamics such as parametric resonance or a phase transition can produce non-Gaussian density perturbations. We show how these perturbations can be calculated using non-linear lattice field theory simulations and the separate universe approximation. In the massless preheating model, we find that some parameter values are excluded while others lead to acceptable but observable levels of non-Gaussianity. This shows that preheating can be an important factor in assessing the viability of inflationary models. | 7 | 10 | 0710.4133 |
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0710 | 0710.4075_arXiv.txt | % Motivated by the emergence of multicore architectures, and the reality that parallelism is rarely used for analysis in observational astronomy, we demonstrate how general users may employ tightly-coupled multiprocessors in scriptable research calculations while requiring no special knowledge of parallel programming. Our method rests on the observation that much of the appeal of high-level vectorized languages like IDL or MatLab stems from relatively simple internal loops over regular array structures, and that these loops are highly amenable to automatic parallelization with OpenMP. We discuss how ISIS, an open-source astrophysical analysis system embedding the \slang\ numerical language, was easily adapted to exploit this pattern. Drawing from a common astrophysical problem, model fitting, we present beneficial speedups for several machine and compiler configurations. These results complement our previous efforts with PVM, and together lead us to believe that ISIS is the only general purpose spectroscopy system in which such a range of parallelism -- from single processors on multiple machines to multiple processors on single machines -- has been demonstrated. | As noted in Noble et al (2006), parallel computation is barely used in astronomical analysis. For example, models in XSPEC (Arnaud 1996), the de facto standard X-ray spectral analysis tool, still run serially on my dual-CPU desktop. In this situation scientists tend to either turn away from models which are expensive to compute or just accept that they will run slowly. Analysis systems which do not embrace parallelism can process at most the workload of only 1 CPU, resulting in a dramatic $1/n$ underutilization of resources as more CPU cores are added. At the same time, however, astronomers are well versed in scripting, particularly with very high-level, array-oriented numerical packages like IDL, PDL, and \slang, to name a few. They combine easy manipulation of mathematical structures of arbitrary dimension with most of the performance of compiled code, with the latter due largely to moving array traversals from the interpreted layer into lower-level code like this C fragment {\small \begin{verbatim} case SLANG_TIMES: for (n = 0; n $<$ na; n++) c[n] = a[n] * b[n]; \end{verbatim} } \noindent which provides vectorized multiplication in \slang. This suggests that much of the strength and appeal of numerical scripting stems from relatively simple loops over regular structures. Having such loops in lower-level compiled codes also makes them ripe for parallelization with \omp\ on shared memory multiprocessors. Proponents contend that conceptual simplicity makes \omp\ more approachable than other parallel programming models, e.g. message-passing in MPI or PVM, and emphasize the added benefit of allowing single bodies of code to be used for both serial and parallel execution. For instance, preceding the above loop with \verb+#pragma omp parallel for+ parallelizes the \slang\ multiplication operator; the pragma is simply ignored by a non-conforming compiler, resulting in a sequential program. | We are witnessing the arrival of serious multiprocessing capability on the desktop: multicore chip designs are making it possible for general users to access many processors. At the granularity of the operating system it will be relatively easy to make use of these extra cores, say by assigning whole programs to separate CPUs. As noted with increasing frequency of late, though, it is not as straightforward to exploit this concurrency {\em within} individual desktop applications. In this paper we demonstrated how we have helped our research colleagues prepare for this eventuality. We have enhanced the vectorization capabilities of \slirp, a module generator for the \slang\ numerical scripting language, so that wrappers may be annotated for automatic parallelization with \omp. This lets \slang\ intrinsic functions be replaced with parallelized versions, at runtime, without modifying a single line of internal \slang\ source. We have shown how \slang\ operators may also be parallelized with relative ease, by identifying key loops within the interpreter source, tagging them with \omp\ directives and recompiling. These simple adaptations, which did not require any changes to the \isis\ architecture or codebase, have yielded beneficial speedups for computations actively used in astrophysical research, and allow the same numerical scripts to be used for both serial and parallel execution -- minimizing two traditional barriers to the use of parallelism by non-specialists: learning how to program for concurrency and recasting sequential algorithms in parallel form. By transparently using \omp\ to effect greater multiprocessor utilization we gain the freedom to explore on the desktop more challenging problems that other researchers might avoid for their prohibitive cost of computation. The \omp\ support now available in GCC makes the techniques espoused here a viable option for many open source numerical packages, opening the door to wider adoption of parallel computing by general practitioners. \begin{figure*}[t] \centering \includegraphics[angle=-90,scale=0.22]{P10.5_3.eps} \hspace*{6mm} \includegraphics[angle=-90,scale=0.22]{P10.5_4.eps} \caption{Aggregate speedup of the Weibull fit function due to the parallelized operators and functions discussed above. Left: \lint, with inflection point at 1907 elements. Right: \solf, with inflection point at 384 elements. } \label{P10.5-fig-2} \end{figure*} | 7 | 10 | 0710.4075 |
0710 | 0710.1650_arXiv.txt | We report on an investigation of the SBS 1520+530 gravitational lens system. We have used archival HST imaging, Keck spectroscopic data, and Keck adaptive-optics imaging to study the lensing galaxy and its environment. The AO imaging has allowed us to fix the lens galaxy properties with a high degree of accuracy when performing the lens modeling, and the data indicate that the lens has an elliptical morphology and perhaps a disk. The new spectroscopic data suggest that previous determinations of the lens redshift may be incorrect, and we report an updated, though inconclusive, value $z_{lens} = 0.761$. We have also spectroscopically confirmed the existence of several galaxy groups at approximately the redshift of the lens system. We create new models of the lens system that explicitly account for the environment of the lens, and we also include improved constraints on the lensing galaxy from our adaptive-optics imaging. Lens models created with these new data can be well-fit with a steeper than isothermal mass slope ($\alpha = 2.29$, where $\rho \propto r^{-\alpha}$) if $H_0$ is fixed at 72\hunit; isothermal models require $H_0 \sim 50$\hunit. The steepened profile may indicate that the lens is in a transient perturbed state caused by interactions with a nearby galaxy. | The strong gravitational lens system SBS 1520+530 (hereafter SBS1520) was first investigated by \citet{chavushyan}, who found that the system consists of a pair of images of a broad absorption line quasar ($z_{src} = 1.855$) separated by 1\farcs6. The lensing galaxy was soon detected with adaptive optics (AO) imaging \citep{crampton} and was assumed to be at the redshift of one of two absorption line systems seen in the spectra of the quasar images. \citet{burud} attempted to deconvolve the lens spectrum from the quasar spectra and found the redshift of the lens to be $z_{lens} = 0.717$. This redshift is broadly consistent with a photometric redshift determined by \citet{faure}. Furthermore, the lens was found to lie along the line of sight to a photometrically identified cluster of galaxies that is expected to be at approximately the same redshift as the lens \citep{burud,faure}. Optical monitoring campaigns have led to the measurement of a time delay between the quasar images of $\sim 130$ days \citep{burud,khamitov}. This time delay provides an additional constraint for determining the mass slope of the lens galaxy \citep[e.g.,][]{rusin} and allows a value of the Hubble Constant to be determined for the system \citep[$H_0 = 51$\hunit assuming an isothermal mass profile;][]{burud}. Note, however, that the mass slope and, thus, the value of $H_0$ depend on the environment surrounding the lens system \citep[e.g.,][]{dobke,auger}. An incorrect understanding of the mass distribution and environment of the lens might account for the anomalously low value of $H_0$ obtained for SBS1520 compared to other lens systems \citep[e.g.,][]{koopmans03,york} and the WMAP \citep[][]{spergel} and \em{Hubble Space Telescope} (\em{HST}) Key Project \citep[][]{freedman} results. In this paper we present new Keck AO and archival \emph{HST} imaging of the lens system that indicates the lensing galaxy may have a disk component. We also present a spectroscopic investigation of the lens environment and provide a new analysis of the lensing galaxy's spectrum which results in an updated lens redshift of $z_{lens} = 0.761$. We discuss the implications of these new observational data on previous analyses performed with SBS1520 and suggest that, in spite of some complexity, this lens provides an interesting platform to investigate dark matter interactions in dense environments. | We have obtained deep AO imaging and optical spectroscopy of the time-delay lens SBS1520. The AO imaging has allowed us to fix the lens galaxy properties with a high degree of precision when performing the lens modeling, and the data indicate that the lens has an elliptical morphology and perhaps a disk. The new spectroscopic data suggest that previous determinations of the lens redshift may be incorrect, and the data also allow us to quantify the lensing contribution of several groups in the immediate foreground and background of the lens. Lens models created with these new data can be well-fit with a steeper than isothermal mass slope ($\alpha = 2.29$) if $H_0$ is fixed at 72\hunit; isothermal models require $H_0 \sim 50$\hunit. \citet{dobke} found that galaxies in overdense environments might have steeper than isothermal mass slopes caused by interactions with other galaxies \citep[also see][]{augerslacs}. This suggests an interpretation that we are observing transient steepening of the mass profile due to galaxy-galaxy interactions and indicates that other lens systems that have obtained anomalously low values of $H_0$ may lie in overdense regions and near an interacting galaxy \citep[e.g., B1600+434;][]{koopmans1600,auger}. Alternatively, SBS1520 can be modeled in a manner consistent with an isothermal profile and $H_0 = 64$\hunit~if if the lens is modeled by a pixelated mass distribution and jointly modeled with other lens systems \citep{read}. These models indicate that twisting ellipticity, triaxial projection effects, or other shape degeneracies may be effecting the parametric analyses of SBS1520 \citep{saha}. However, there are still several ambiguities in the data that need to be resolved before making definitive claims about the profile of SBS1520, particularly in the context of the interaction-induced steepening scenario. While we have argued that the lens redshift is likely to be $z = 0.761$ and not $z = 0.717$, the data are not conclusive. Furthermore, our modeling has assumed that all of the neighbor galaxies are at the group redshift; if this is the case, the $z = 0.76$ group centroid would be pulled closer to the lens and the group would therefore provide a larger contribution to the lens model. If this is not the case, the neighboring galaxies might have a smaller impact on the lens model. This is particularly important for Galaxy M, as this is the neighboring galaxy that most affects the lens model but also has colors least like the lens galaxy compared to the other field galaxies. It is also important to verify that at least one of the neighboring galaxies is at the same redshift as the lens because this is a requirement of the interaction-driven steepening hypothesis. Finally, obtaining a dynamical estimate of the lens mass would help to further constrain models and potentially distinguish between shape degeneracies and the mass slope degeneracy. | 7 | 10 | 0710.1650 |
0710 | 0710.5073_arXiv.txt | Due to the non-commutation of spatial averaging and temporal evolution, inhomogeneities and anisotropies (cosmic structures) influence the evolution of the averaged Universe via the cosmological backreaction mechanism. We study the backreaction effect as a function of averaging scale in a perturbative approach up to higher orders. We calculate the hierarchy of the critical scales, at which $10\%$ effects show up from averaging at different orders. The dominant contribution comes from the averaged spatial curvature, observable up to scales of $\sim 200~\mbox{Mpc}$. The cosmic variance of the local Hubble rate is $10\%$ $(5\%)$ for spherical regions of radius 40 $(60)~\mbox{Mpc}$. We compare our result to the one from Newtonian cosmology and Hubble Space Telescope Key Project data. | 7 | 10 | 0710.5073 |
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0710 | 0710.5303_arXiv.txt | We have updated predictions for high energy neutrino and antineutrino charged current cross-sections within the conventional DGLAP formalism of NLO QCD using a modern PDF fit to HERA data, which also accounts in a systematic way for PDF uncertainties deriving from both model uncertainties and from the experimental uncertainties of the input data sets. Furthermore the PDFs are determined using an improved treatment of heavy quark thresholds. A measurement of the neutrino cross-section much below these predictions would signal the need for extension of the conventional formalism as in BFKL resummation, or even gluon recombination effects as in the colour glass condensate model. | \label{sec:intro} Predictions of neutrino cross-sections at high energies have sizeable uncertainties which derive largely from the measurement uncertainties on the parton distribution functions (PDFs) of the nucleon. In the framework of the quark-parton model, high energy scattering accesses very large values of $Q^2$, the invariant mass of the exchanged vector boson, and very small values of Bjorken $x$, the fraction of the momentum of the incoming nucleon taken by the struck quark. Thus when evaluating uncertainties on high energy neutrino cross-sections it is important to use the most up to date information from the experiments at HERA, which have accessed the lowest-$x$ and highest $Q^2$ scales to date. The present paper uses the formalism of the ZEUS-S global PDF fits~\cite{Chekanov:2002pv}, updated to include {\em all} the HERA-I data. Conventional PDF fits use the Next-to-leading-order (NLO) Dokshitzer-Gribov-Lipatov-Altarelli-Parisi (DGLAP) formalism~\cite{Altarelli:1977zs,Gribov:1972ri,Lipatov:1974qm,Dokshitzer:1977sg} of QCD to make predictions for deep inelastic scattering (DIS) cross-sections of leptons on hadrons. At low-$x$ where the gluon density is rising rapidly it is probably necessary to go beyond the DGLAP formalism in order to sum $\ln(1/x)$ diagrams, as in the Balitsky-Fadin-Kuraev-Lipatov (BFKL) formalism~\cite{Kuraev:1977fs,Balitsky:1978ic,Lipatov:1985uk} (for recent work see~\cite{White:2006xv,Altarelli:2005ni,Ciafaloni:2006yk}). An alternative approach is to consider non-linear terms which describe gluon recombination as in the colour glass condensate model~\cite{Iancu:2003xm} which has had considerable success in explaining RHIC data~\cite{JalilianMarian:2005jf}. A recent suggestion is to use a structure function consistent with HERA data that saturates the Froissart unitarity bound and thus predicts a $\ln^2 s$ dependence of the cross-section~\cite{Berger:2007ic}. Such approaches are beyond the scope of the present paper, which is concerned with the more modest goal of estimating the uncertainties on high energy neutrino cross-sections which are compatible with the conventional NLO DGLAP formalism. The motivation is to provide an update on the neutrino cross-sections in the literature \cite{Gandhi:1998ri} which are widely used e.g. for estimating event rates in neutrino telescopes such as Baikal \cite{Antipin:2007zz}, ANTARES \cite{Aslanides:1999vq} and IceCube \cite{icecube}, cosmic ray observatories such as HiRes \cite{Martens:2007ff} and Auger \cite{auger}, and radio detectors such as GLUE \cite{Gorham:2003da}, FORTE \cite{Lehtinen:2003xv}, RICE \cite{Kravchenko:2002mm} and ANITA \cite{Barwick:2005hn}. As a corollary, if cross-sections much outside these limits are observed, it would be a clear signal of the need for extensions to conventional formalism. To date no unambiguous signals which require such extensions have been observed. The prospect for measuring the cross-section using very high energy cosmic neutrinos in order to distinguish between theoretical suggestions for gluon dynamics at low $x$ has been discussed by us elsewhere \cite{Anchordoqui:2006ta}. Previous work on estimating high energy neutrino cross-sections~\cite{Gandhi:1998ri} used PDF sets which no longer fit modern data from HERA~\cite{Tung:2004rw} and an {\em ad hoc} procedure for estimating PDF uncertainties. The present work improves on this in several respects. Firstly, we use a recent PDF analysis which includes data from all HERA-I running~\cite{Chekanov:2002pv}. Secondly, we take a consistent approach to PDF uncertainties --- both model uncertainties and, more importantly, the uncertainties which derive from the correlated systematic errors of the input data sets \cite{CooperSarkar:2002yx}. Thirdly, we use NLO rather than LO calculations throughout. Fourthly, we use a general-mass variable flavour number scheme~\cite{Thorne:1997ga,Thorne:2006qt} to treat heavy quark thresholds. | \label{sec:conc} We have calculated the charged current neutrino cross-section at NLO in the Standard Model using the best available DIS data along with a careful estimate of the associated uncertainties. As mentioned earlier, there are further uncertainties associated with QCD effects at very low $x$ which are not addressed in the DGLAP formalism. When $x$ is sufficiently small that $\alpha_\mathrm{s}\,\ln (1/x) \sim 1$, it is necessary to resum these large logarithms using the BFKL formalism. Whereas such calculations at leading-log suggest an even steeper rise of the gluon structure function at low $x$ (which would imply a higher $\nu-N$ cross-section), this rise is not so dramatic at next-to-leading-log; for a recent application of NLL BFKL resummation to deep inelastic scattering see~\cite{White:2006yh}. Moreover both the DGLAP and the BFKL formalisms neglect non-linear screening effects due to gluon recombination which may lead to saturation of the gluon structure function. This has been modelled in the colour dipole framework in which DIS at low $x$ is viewed as the interaction of the $q \bar q$ dipole to which the gauge bosons fluctuate. An unified BFKL/DGLAP calculation \cite{Kwiecinski:1998yf} supplemented by estimates of screening and nuclear shadowing effects, predicts a {\em decrease} of the $\nu-N$ cross-section by $20-100\%$ at very high energies $E_\nu \sim 10^8-10^{12}$~GeV~\cite{Kutak:2003bd}. An alternative approach uses the colour glass condensate formalism~\cite{Iancu:2003xm} and predicts a similar suppression when a dipole model~\cite{Kharzeev:2004yx} which fits data from RHIC is used~\cite{Henley:2005ms}. The predicted cross-section is even lower~\cite{Henley:2005ms} if a different dipole model~\cite{Bartels:2002cj} developed to fit the HERA data is used and the gluon distribution is assumed to decrease for $x < 10^{-5}$. Other possibilities for the behaviour of the high energy $\nu-N$ cross-section have also been discussed~\cite{Berger:2007ic,Jalilian-Marian:2003wf}. Detectors for UHE cosmic neutrinos would be able to probe such new physics if they can establish deviations from the perturbative DGLAP prediction. Hence we recommend our calculated values for estimation of the baseline event rates in neutrino telescopes and for use in event generators such as ANIS \cite{Gazizov:2004va}. While the expected neutrino fluxes (e.g. from the sources of the observed high energy cosmic rays) are rather uncertain, experiments can in principle exploit the different dependence on the cross-section of the rate of Earth-skimming and quasi-horizontal events \cite{Anchordoqui:2006ta}. | 7 | 10 | 0710.5303 |
0710 | 0710.5629_arXiv.txt | The specifications of the Atacama Large Millimeter Array (ALMA) have placed stringent requirements on the mechanical performance of its antennas. As part of the evaluation process of the VertexRSI and Alcatel EIE Consortium (AEC) ALMA prototype antennas, measurements of the path length, thermal, and azimuth bearing performance were made under a variety of weather conditions and observing modes. The results of mechanical measurements, reported here, are compared to the antenna specifications. | \PARstart{T}{he} Atacama Large Millimeter Array (ALMA) for astronomical observations at millimeter and sub--millimeter wavelengths (up to the Terra--Hertz region) needs antennas of high mechanical precision and of understandable and predictable behaviour. This behaviour must be established for structural deformations due to gravity, temperature changes, and wind loads. This means, in particular, that a high reflector surface precision, pointing and phase stability must be maintained under all motions of tracking and mapping. We present a summary of tests of the mechanical and thermal behaviour of the 12m diameter VertexRSI and AEC ALMA prototype antennas, built at the VLA site (2000m altitude), New Mexico, USA. The tests were made at several intervals between March 2003 and April 2004, and concentrated primarily on the verification of the antenna specifications, of path length variations and parameters which influence the pointing. The data were also analyzed to understand the general behaviour of the antennas. In the investigation we have paid attention to the fact that variations in the behaviour of the antennas may be predictable or sporadic. We believe that repeatable and/or predictable variations can to a large extent be considered in the pointing model, or any other correction device. The antennas were tested in stationary position, under sidereal tracking, On--the--Fly (OTF) mapping, and in Fast--Switching mode (FSW). A large amount of data was collected during commissioning and thus refer to all types of tracking, OTF, FSW, and unintended 'shaking'. A more extended summary of these test results was reported by the Antenna Evaluation Group to the National Radio Astronomy Observatory (NRAO) and European Southern Observatory (ESO) (which forms the ALMA partnership) for selection of the ALMA production antenna(s). An overview of these performance results was presented in \cite{Mangum2006}. The current paper presents a more detailed analysis of the mechanical performance measurements made of the ALMA prototype antennas. | The measurements indicate that the path length specifications are fulfilled on both antennas, at least during time intervals of 1/2 to 1 hour. The path length variation is primarily due to unavoidable residual thermal dilatation of the (insulated) antenna steel components, and may span $\sim$\,200\,$\mu$m within a day. The path length variation can be predicted with high precision from temperature measurements at a few positions of the steel components, either used in empirical relations or the finite element model. It is expected that identical antennas will experience similar temperature variations of the ambient air so that the differential effect may even be smaller than stated here. Wind at speeds below the specification limit (9\,m/s), and OTF and FSW motions of the antennas, do not degrade the path length stability. As far as possible to measure, the antennas show similar behaviour of damping of the thermal environment, \ie\ the ambient air temperature and the solar radiation. The BUS of the VertexRSI antenna shows a good temperature homogeneity, even under full exposure to Sun shine. Altough the AZ bearings have a higher order azimuth dependent wobble, the effect can be considered in the pointing model with an accuracy better than 0.6 arcsecond. On the VertexRSI prototype antenna, the wobble was very stable with time. | 7 | 10 | 0710.5629 |
0710 | 0710.2009_arXiv.txt | \label{sec:intro} There has been much interest in recent years in cosmological applications of string theory. The availability of precision data relating to the cosmic microwave background (CMB) opens up the possibility of constraining the phenomenology of string--inspired models, in particular regarding the dynamics of the inflationary epoch. Essentially, signatures of the extreme conditions governing the behaviour of the early universe are amplified by inflationary expansion and can be accessible by modern day observations. The typical energy scales involved far exceed those obtainable in particle collider experiments. In this paper we consider a particular inflation scenario, the Dirac--Born--Infeld (DBI) model of~\cite{Silverstein:2003hf,Alishahiha:2004eh}. In the usual formulation of this setup, a D3 brane rolls down a warped throat towards a $\overline{\text{D3}}$ brane where the interbrane separation is identified with the inflaton field. The brane moves relativistically, but its speed is curtailed by the warped geometry so that the potential energy dominates and inflation can occur. The observable consequences of this model depend of course upon the choice of solution for the warped throat. The original DBI proposal used a AdS--like geometry, with an artificially imposed cutoff where the throat joins the bulk geometry, which in principle is unknown although irrelevant for cosmological implications. Such a geometry is unstable, and more properly one should consider a solution of the supergravity field equations with non--trivial D--brane fluxes to stabilise the compact geometry by dynamically generating a cut--off. Not many such solutions are known analytically, but two examples are the Klebanov--Strassler~\cite{Klebanov:2000hb} (KS) and Maldacena--Nu{\~n}ez~\cite{Maldacena:2000yy} (MN) results.\\ Several investigations of DBI inflation from the KS geometry abound in the literature~\cite{Kecskemeti:2006cg,Shiu:2006kj,Easson:2007dh,Shandera:2006ax,Bean:2007hc}, including systematic comparisons to current cosmological data. In this paper we show that it is also possible to examine the inflationary consequences of the BGMPZ geometries introduced in~\cite{Butti:2004pk}. These are a one--parameter family of geometries describing a deformation of the KS solution\footnote{The BGMPZ solutions originally arose in a different context, that of the AdS / CFT correspondence, as the baryonic branch of the KS solution.}, and interpolating smoothly between the KS and MN backgrounds. Although the metric is not known fully analytically, one can solve for it numerically and thus obtain results for numerous observables. The hope is that these can be compared with experimental data for quantities such as spectral indices, which may then allow one to extract geometrical information about the throat--geometry. To this aim, we show examples of how these quantities vary as a function of the BGMPZ parameter $\xi$ which characterises the family of solutions, at typical points in parameter space. We note that inflationary consequences of these geometries were also considered in~\cite{Dymarsky:2005xt}, for slow-roll brane inflation rather than the DBI setup considered in this paper.\\ The structure of the paper is as follows. In section \ref{sec:dbi} we summarise the salient details of DBI inflation relevant for our purposes, followed by a brief explanation of BGMPZ backgrounds. In section \ref{sec:calc} we explain the details of the numerical procedure adopted in solving the BGMPZ equations. Example results are presented in section \ref{sec:results}, and in section \ref{sec:discuss} we discuss our results before concluding. | \label{sec:discuss} In this paper we have explicitly demonstrated the possibility of calculating cosmological observables in the DBI inflation setup using a one parameter family of type IIB supergravity solutions that describe the geometry of a warped throat, the BGMPZ solutions of Butti et al.~\cite{Butti:2004pk} that interpolate smoothly between the Klebanov--Strassler and the Maldacena--Nu{\~n}ez solution. We have provided examples of cosmological parameters, namely spectral indices, that can be calculated from the underlying geometry.\\ The solution for the metric of the geometries in question is not possible analytically. Therefore numerical methods have been used and shown to provide an adequate representation of the metric in terms of numerical precision. Instabilities in the derivatives entering the equations have been dealt with by matching to known power series expansions of the metric functions at asymptotically small and large values of the radial coordinate.\\ We presented warp factors for two almost extremal ends of the family of solutions parametrised by~$\xi$: the Klebanov--Strassler throat ($\xi=1/2$), and a geometry close to the Maldacena--Nu{\~ n}ez solution ($\xi=0.167$). The qualitative behaviour of the warp factors is seen to be different. Both solutions show a flat warp factor at low values of the rescaled radial coordinate $\phi$, corresponding to a dynamically generated cutoff, and an asymptotically AdS behaviour at large $\phi$ values as expected. However, the warp factor as it moves away in $\xi$ from the KS solution develops a shoulder at intermediate values, whose slope is in general different from the eventual asymptotic behaviour. This can be effectively parameterised by a function of form (\ref{fpar2}), thus explicitly demonstrating that phenomenological fits of DBI inflation to cosmological data that assume a warp factor of this form indeed correspond to exactly realisable warped throat geometries in a known string theory compactification.\\ We presented examples of scalar spectral indices and the ratio of tensor to scalar modes calculated using the different geometries. The different qualitative behaviour observed in the warp factors carries through to the spectral indices, and a quantitative difference between the values of the indices from different solutions at 55 $e$--folds before the end of inflation is potentially measurable. Constraints on non-Gaussianity and measured values of the scalar spectral index could already be used to rule out some regions of $\xi$, when it is considered alongside the rest of the DBI parameter space. Very generically, we have found that the amount of non--Gaussianity increases as one moves away from the KS solution (all other parameters being equal).\\ Since the warping changes qualitatively as one varies $\xi$, it is also conceivable that that certain constraints on the parameter space may be relaxed by including $\xi$ as an additional parameter with respect to existing analyses such as~\cite{Lorenz:2007ze}. This issue certainly deserves further study. \subsection*{Acknowledgements} We would like to thank Gary Shiu for initiating this project and for very helpful correspondence. We acknowledge informative discussions with Marieke Postma. Our research is supported by the Dutch Foundation for Fundamental Research of Matter (FOM). | 7 | 10 | 0710.2009 |
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0710 | 0710.3112_arXiv.txt | We study CP-violation effects when neutrinos are present in dense matter, such as outside the proto-neutron star formed in a core-collapse supernova. Using general arguments based on the Standard Model, we confirm that there are no CP-violating effects at the tree level on the electron neutrino and anti-neutrino fluxes in a core-collapse supernova. On the other hand significant effects can be obtained for muon and tau neutrinos even at the tree level. We show that CP violating effects can be present in the supernova electron (anti)neutrino fluxes as well, if muon and tau neutrinos have different fluxes at the neutrinosphere. Such differences could arise due to physics beyond the Standard Model, such as the presence of flavor-changing interactions. | Recent results from solar, atmospheric and reactor experiments have significantly improved our knowledge of the neutrino mass differences and of two of the mixing angles. If the remaining mixing angle, $\theta_{13}$, is relatively large there is a possibility that violation of CP symmetry may be observable in the neutrino sector. Currently planned and future experiments will have improved sensitivities to the value of this angle (see e.g. \cite{Guo:2007ug,Ardellier:2006mn,Kato:2007zz}). Effects of CP-violation in accelerator neutrino oscillation experiments have been extensively investigated \cite{Dick:1999ed,Peres:2003wd,Ishitsuka:2005qi,Barger:1980jm,Karagiorgi:2006jf,Volpe:2006in}. The discovery of a non zero Dirac delta phase might help our understanding of the observed matter-antimatter asymmetry of the universe \cite{Pascoli:2006ie,Pascoli:2006ci,Mohapatra:2006se,Barger:2003gt}. Besides studies on terrestrial experiments with man-made sources, a few recent works have addressed CP-violation with neutrinos from astrophysical sources (see e.g. \cite{Akhmedov:2002zj,Winter:2006ce}). The purpose of the present paper is to explore possible effects coming from the CP-violating phase in dense matter, such as that encountered in core-collapse supernovae. Core-collapse supernovae occur following the stages of nuclear burning during stellar evolution after an iron core is formed. The iron cores formed during the evolution of massive stars are supported by the electron degeneracy pressure and hence are unstable against a collapse during which most of the matter is neutronized. Once the density exceeds the nuclear density this collapse is halted. Rebounding pressure waves break out into a shock wave near the sonic point where the density reaches the nuclear density. Evolution of this shock wave and whether it produces an explosion is a point of current investigations. However, it is observed that the newly-formed hot proto-neutron star cools by neutrino emission. Essentially the entire gravitational binding energy of eight or more solar mass star is radiated away in neutrinos. Although the initial collapse is a very orderly (i.e. low entropy) process, during the cooling stage at later times a neutrino-driven wind heats the neutron-rich material to high entropies \cite{Woosley:1992ek,Woosley:1994ux,Takahashi:1994yz}. Neutrino interactions play a very important role in the evolution of core-collapse supernovae and in determining the element abundance \cite{Balantekin:2003ip}. Neutrino heating is a possible mechanism for reheating the stalled shock \cite{Bethe:1984ux}. A good fraction of the heavier nuclei were formed in the rapid neutron capture (r-process) nucleosynthesis scenario \cite{Burbidge:1957vc}. Core-collapse supernovae are one of the possible sites for the r-process nucleosynthesis. A key quantity for determining the r-process yields is the neutron-to-seed nucleus ratio, determined by the neutron-to-proton ratio, which is controlled by the neutrino fluxes. In addition, recent work indicates that neutrino-neutrino interactions plays a potentially very significant role in supernovae \cite{Samuel:1993uw,Sigl:1992fn,Balantekin:2004ug,Duan:2006jv,Balantekin:2006tg,Hannestad:2006nj}. In this paper we study CP violation aspects in the core-collapse supernova environment. We first analyze analytically and in general terms, how the neutrino propagation equations and the evolution operator are modified in matter, in presence of a non-zero Dirac delta phase. We obtain a general formula which is valid for any matter density profile\footnote{Such findings are in agreement with what was found in Ref. \cite{Yokomakura:2002av}.}. In particular we demonstrate that, as in vacuum, the electron (anti)neutrino fluxes are independent of the phase $\delta$, if mu and tau neutrinos have the same fluxes at the neutrinosphere in the supernova\footnote{A remark on this aspect was made in \cite{Yoshida:2006sk}.}. On the other hand the electron (anti)neutrino fluxes will depend on $\delta$, if mu and tau neutrinos have different fluxes at the neutrinosphere, at variance with what was found in \cite{Akhmedov:2002zj}. We present numerical calculations on possible CP violation effects on the mu and tau neutrino fluxes as well as on the electron (anti-)neutrino flux and the electron fraction. The latter can only appear if physics beyond the Standard Model, such as flavor changing interactions, induces differences on the mu and tau neutrino initial total luminosities and/or temperatures. Finally we calculate the effects from the CP violating phase on the number of events in an observatory on earth. The plan of this paper is as follows. In Section II we present the general formalism to describe the neutrino evolution in presence of the $\delta$ phase. The formulas concerning neutrino fluxes and the electron fraction in the supernova environment are recalled in Section III. Numerical results on these quantities are presented in Section IV. Conclusions are made in Section V. | \noindent In this work we have analyzed possible effects induced by the CP violating Dirac phase in a dense environment such as the core-collapse supernovae. Our major result are that in matter: i) significant effects are found on the muon and tau neutrino fluxes for a non-zero CP violating phase; ii) important effects are also found on the electron (anti)neutrino fluxes if the $\nu_{\mu}$ and $\nu_{\tau}$ neutrino fluxes differ at the neutrinosphere. On the other hand the usual assumption of ignoring the CP violating phase made in the literature is justified if contributions from physics beyond the Standard Model is small and the $\nu_{\mu}$ and $\nu_{\tau}$ fluxes are equal at emission. We have calculated the events in an observatory on earth and shown that effects at the level of $5 \%$ are present on the number of events as a function of neutrino energy. Recent calculations have shown that the inclusion of neutrino-neutrino interaction introduces new features in the neutrino propagation in supernovae. A detailed study of the neutrino evolution with the CP violating phase, the neutrino-neutrino interaction as well as loop induced neutrino refractive indices will be the object of further work. | 7 | 10 | 0710.3112 |
0710 | 0710.3748_arXiv.txt | We used multiwavelength data (H{\sc i}, FUV, NUV, R) to search for evidence of star formation in the intragroup medium of the Hickson Compact Group 100. We find that young star-forming regions are located in the intergalactic H{\sc i} clouds of the compact group which extend to over 130 kpc away from the main galaxies. A tidal dwarf galaxy candidate is located in the densest region of the H{\sc i} tail, 61 kpc from the brightest group member and its age is estimated to be only 3.3 Myr. Fifteen other intragroup H{\sc ii} regions and TDG candidates are detected in the GALEX FUV image and within a field 10$'$$\times$10$'$ encompassing the H{\sc i} tail. They have ages $<$200 Myr, H{\sc i} masses of 10$^{9.2-10.4}$ M$_{\odot}$, 0.001$<$ SFR $<$0.01 M$_{\odot}$ yr$^{-1}$, and stellar masses 10$^{4.3}$--10$^{6.5}$ M$_{\odot}$. The H{\sc i} clouds to which many of them are associated have column densities about one order of magnitude lower than N(H{\sc i})$\sim$10$^{21}$ cm$^{-2}$. | The environment of galaxies play an important role on determining their overall properties. One of the most interesting environmental effects is seen in interacting systems which contain stripped H{\sc i} gas in the intragroup/intergalactic medium, instead of around galaxies. Three independent studies, Mendes de Oliveira et al. (2004, for the Stephan's Quintet), Mendes de Oliveira et al. (2006, for HCG 31), Ryan-Weber et al. (2004, for NGC 1533, HCG 16, ESO 149-G003) and Oosterloo et al. (2004, for the loose group around NGC 1490) have surveyed systems with H{\sc i} intergalactic clouds and have shown that these are associated with actively star forming regions, the so-called {\it intergalactic H{\sc ii} regions}. These objects seem to be similar to the H{\sc ii} regions in our Milky Way, but are located in the intragroup medium and have high metallicities (close to solar in most cases). The fate of these types of objects and their importance in galaxy formation and evolution, enrichment of the intergalactic medium and globular cluster formation is still debatable. In order to better address these issues, we embarked on a multiwavelength analysis of the environment of interacting systems and present here the results based on FUV, NUV, optical, and H{\sc i} data of a Hickson compact group, HCG100. The last group in Hickson's catalog of compact groups of galaxies (Hickson 1982) is at v$_{\rm R}$ = 5336 km/s (z=0.0178, Hickson et al. 1992) and it is formed by four late-type galaxies with accordant redshifts: a bright central Sb galaxy (HCG100a), an irregular galaxy with an optical tidal tail (HCG100b), a late-type barred spiral (HCG100c) with central H$\alpha$ emission (Vilchez \& Iglesias-P\'aramo 1998) and a late-type edge-on spiral (HCG100d). HCG100a, b and c show strong evidence of interaction, demonstrated not only by peculiarities in their morphologies but also in their velocity fields (Plana et al. 2003). H100b shows a tidal tail connecting with a faint galaxy not originally classified as a member of the group by Hickson. Past encounters are also confirmed by an extended H{\sc i} tail (Verdes-Montenegro et al. 2006). Therefore, the well-advanced dynamical state of HCG100 makes it an ideal target for searching for intergalactic H{\sc ii} regions that might have been triggered by tidal interaction. This paper is organized as follows: \S 2 describes the data, \S 3 discusses the age, masses and star formation rate estimates, \S 3.1 and \S 3.2 compare the intragroup regions in HCG100 with other intergalactic H{\sc ii} regions and high$-z$ UV-luminous galaxies. Finally, \S4 summarizes the main conclusions. Throughout this paper, we use a cosmology with $\Omega_{\rm M}=0.3$, $\Omega_{\Lambda}=0.7$~and $h=0.7$. Magnitudes are given in the AB-system and the adopted distance to HCG100 is 76.3 Mpc. \section {The data} HCG100 was observed with the Galaxy Evolution Explorer (GALEX) mission in the far and near ultraviolet (FUV $\lambda$$_{\rm eff}$=1528\AA, NUV $\lambda$$_{\rm eff}$=2271\AA) as part of our Cycle 1 program (GI\#73) to observe compact groups of galaxies. In Fig. \ref{fuvhi_10arcmin} and Fig.\ref{nuvhi_10arcmin} we show a cutout of a 10$'$$\times$ 10$'$ window of the FUV and NUV images (GALEX field of view is 1$^{\circ}$.28 and 1$^{\circ}$.24 in FUV and NUV, respectively and pixel scale is 1.5 arcsec pixel$^{-1}$) together with the distribution of atomic hydrogen H{\sc i} (PI Verdes-Montenegro). We have also obtained an R-band image with the CTIO Blanco 4m telescope and a mosaic II CCD imager. Three 300s images of the group were taken, at a median seeing of 1.1$''$. Each frame covered a 40 arcmin field on a side, at a pixel scale of 0.27$''$/pixel, but we only used a 10 arcmin field around the object, given that this was the area of interest. The data was bias subtracted and flat fielded using standard procedures (using the package mscred in IRAF\footnote{IRAF is distributed by the National Optical Astronomy Observatories, which are operated by the Association of Universities for Research in Astronomy, Inc., under cooperative agreement with the National Science Foundation.}). A flatfield was constructed from a combination of dark sky frames and twilight flats which worked well, given that the background, after flatfielding was done, showed an rms variation no greater than 1\% over the whole field. The night was not photometric and therefore no calibration stars were observed. Instead, we chose to use two galaxies previously observed by Rubin et al. (1991), HCG 100c and HCG 100d, for obtaining the zero point of the photometry (see below). \subsection{Flux Calibration} FUV and NUV fluxes were calculated using Morrissey et al. (2005) m$_{\lambda}$=-2.5 log[F$_{\lambda}$/a$_{\lambda}$] + b$_{\lambda}$, where a$_{FUV}$ = 1.4 $\times$ 10$^{-15}$ erg s$^{-1}$ cm$^{-2}$ \AA$^{-1}$, a$_{NUV}$=2.06$\times$ 10$^{-16}$ erg s$^{-1}$ cm$^{-2}$ \AA$^{-1}$, b$_{FUV}$=18.82 and b$_{NUV}$=20.08 for FUV and NUV, respectively. Fluxes were multiplied by the effective filter bandpass ($\Delta$$\lambda$$_{FUV}$=269\AA\, and $\Delta$$\lambda$$_{NUV}$=616\AA) to give units of erg s$^{-1}$ cm$^{-2}$ and luminosities were calculated for a distance D=76.3 Mpc. The R-band image was calibrated to match Rubin et al. (1991) magnitudes for galaxies HCG100c and HCG100d. Total integrated R magnitudes of these objects are given in their Table 3 (no errors were quoted), and these were compared with the total instrumental magnitudes we obtained for these galaxies using two methods: (1) through aperture photometry in IRAF (using the task daophot.phot) and (2) using the program SExtractor (Bertin and Arnouts et al. 1996). The zero points obtained for each of these objects, and for each method, agreed within 0.1 mag, and we used an average of the values as our final zero point. The lack of a proper calibration, does not significantly change our results as shown in \S~3. The H{\sc i} data taken with the VLA (configuration D, PI Verdes-Montenegro) is not able to resolve the intragroup objects since the beamwidth is 61.0$''$$\times$55.23$''$ (22.6~kpc $\times$ 20.4~kpc at 76.3~Mpc) and objects have sizes $<$16$''$. Therefore, our measurements of the H{\sc i} fluxes were centered on the FUV detections but measured within one beam. Fluxes were converted into mass by using the relation: \begin{equation} M_{HI}=2.356 \times 10^5 F_{HI} D^2 \end{equation} where $D$ is the distance to the group in Mpc, $F_{HI}$ is the H{\sc i} flux in Jy km/s. Column Densities were calculated by applying the following relation \begin{equation} N_{HI}=1.82\times 10^{18} (\lambda ^2 / (2.65 \Theta ^2) F \end{equation} where $\Theta$=beamwidth in arcmin$^{2}$, $\lambda$ is the wavelength (21cm) and F is the H{\sc i} flux density in Jy/beam km/s (Rohlfs \& Wilson 2000). \subsection{Source Detection} We used SExtractor (SE, Bertin \& Arnouts 1996) to detect sources in the FUV and matched that catalog with the NUV and R-band catalogs. Therefore, only objects with FUV detections were included in our final catalog. In this paper we concentrate on all objects which were detected in a region of 10$'$$\times$ 10$'$ centered on the H{\sc i} tail (Verdes-Montenegro et al. 2006). We used SE's Kron elliptical apertures to measure magnitudes (mag$_{-}$auto) in FUV, NUV and R-bands. Although mag$_{-}$auto is often used to measure total magnitudes in galaxy surveys (e.g. Bell et al. 2004, de Mello et al. 2006, Zucca et al. 2006), high uncertainties might be expected at the faint magnitudes due to the assumption that the sky background has Gaussian random noise without source confusion (Brown et al. 2007). However, since we are comparing data taken with different resolutions and the fact that UV light does not necessarily peak at the same coordinate as the optical light, mag$_{-}$auto performs better than the other SExtractor choices, as long as a careful match between the different catalogs is performed. We matched the coordinates of the sources in the 3-bands within 1$''$ to 5$''$ and after visually checking each identification we chose 3$''$ as our best matched catalog with FUV, NUV and R-detections. The cross-match in the three bands resulted not only in the four most luminous members of HCG 100 (Table 1) but also 16 other objects, within the 10$'$$\times$10$'$ field (Table 2). Nine of these sources are within the H{\sc i} tail: \# 3 is the most distant object from the group (381 arcsec or 137.2 kpc from HCG100a) and is located in the southern tip of the H{\sc i} tail, \#4 is far from all of the brightest members of the group and located in the densest H{\sc i} region, \#5 and \#6 are close to HCG100c, \#8 is isolated and 150 arcsec (53.9 kpc) from HCG100a, \#9 is close to HCG100a, \#13 is a small galaxy not originally classified as a member of the HCG100 and located in the optical tidal tail of HCG100b, \#14 and \#15 are also close to the tidal tail. The other six objects have no H{\sc i} detection but are still within the chosen field. SE's magnitudes in FUV, NUV and R (mag$_{-}$auto) were corrected for foreground Galactic extinction using E(B-V)=0.081 and A$_{R}$=E(B-V)$\times$2.634 (Schlegel et al. 1998), A$_{FUV}$=E(B-V)$\times$8.29 and A$_{NUV}$=E(B-V)$\times$8.18 (Seibert et al. 2005). In Table 3 we list the H{\sc i} mass per beam in the vicinity of nine of the intragroup objects detected, the other 7 objects are below the detection limit ($<1.7 \times 10^{8}$). However, it is important to note that, due to the low resolution of the data, the HI masses of the objects near large galaxies are contaminated by the HI masses of the latter. Only objects \#3, \#4 and \#8 are far away from bright members of the group and this contamination could be avoided. Objects \#13, \#14, \#15 besides being close to H100b are within the same beam, therefore the HI masses listed in Table 3 are not the individual masses of each object but the total mass within one beam centered in that region. As seen in Fig.\ref{fuvhi_10arcmin} most of the intragroup objects are located in the outskirts of the H{\sc i} contours where column densities range from 7.5 $\times 10^{19}$ to 5$\times 10^{20}$ cm$^{-2}$, except for object \#4 which is located in a peak of H{\sc i} where NH{\sc i} is 1.2$\times 10^{21}$ cm$^{-2}$. Therefore, H{\sc i} clouds to which many of them are associated have column densities about one order of magnitude lower than the N(HI)$\sim$10$^{21}$ cm$^{-2}$, value thought to be required for triggering star formation (e.g. Skillman et al. 1988). Due to the low resolution of the HI data, all values are lower limits to the true HI column density. \begin{deluxetable}{cccccccc} \tablecaption{HCG100 Properties} \tablewidth{0pt} \tablehead{ \colhead{ID} & \colhead{Morphology$^{\rm a}$} & \colhead{Velocity$^{\rm b}$} & \colhead{FUV$^{\rm c}$} & \colhead{NUV} & \colhead{R$^{\rm d}$} & \colhead{FUV$_{\rm corr}$$^{\rm e}$} & \colhead{NUV$_{\rm corr}$} } \startdata HCG100 a & Sb & 5323 &17.30 & 16.78 & 12.5 & 16.12 & 16.63\\ HCG100 b & Sm & 5163 &17.49 & 17.10 & 14.1 & 16.44 & 16.82\\ HCG100 c & SBc& 5418 &18.46 & 17.91 & 14.7 & 17.25 & 17.79\\ HCG100 d & Scd& &19.08 & 18.65 & 15.5 & 17.99 & 18.41\\ \enddata \tablenotetext{a}{Morphology from Plana et al. 2003} \tablenotetext{b}{Systemic Velocity in kms$^{-1}$ from Plana et al. 2003} \tablenotetext{c}{FUV and NUV magnitudes were obtained using IRAF task ellipse.} \tablenotetext{d}{R-band magnitudes are from Rubin et al. 1991.} \tablenotetext{e}{Extinction corrections using Seibert et al. (2005) for FUV and NUV.} \end{deluxetable} \begin{deluxetable}{ccccrrrr} \tabletypesize{\scriptsize} \tablecaption{FUV sources within HCG100 $10' \times 10'$ field} \tablewidth{0pt} \tablehead{ \colhead{ID} & \colhead{RA} & \colhead{Dec} & \colhead{R$^{\rm a}$} & \colhead{NUV} & \colhead{FUV} & \colhead{FUV-R} & \colhead{FUV-NUV} } \startdata 1 & 0.2410& 13.0647& 18.80 $\pm$ 0.01& 20.56 $\pm$ 0.04& 20.44 $\pm$ 0.06& 1.64 $\pm$ 0.06 & -0.12 $\pm$ 0.07\\ 2 & 0.2526& 13.1161& 19.54 $\pm$ 0.01& 21.61 $\pm$ 0.08& 21.48 $\pm$ 0.10& 1.94 $\pm$ 0.10 & -0.13 $\pm$ 0.13\\ 3 & 0.2722& 13.0236& 19.09 $\pm$ 0.01& 21.61 $\pm$ 0.08& 21.21 $\pm$ 0.09& 2.12 $\pm$ 0.09 & -0.40 $\pm$ 0.12\\ 4 & 0.2929& 13.0859& 20.25 $\pm$ 0.04& 21.26 $\pm$ 0.07& 21.09 $\pm$ 0.09& 0.84 $\pm$ 0.10 & -0.17 $\pm$ 0.12\\ 5 & 0.2933& 13.1377& 19.42 $\pm$ 0.01& 22.43 $\pm$ 0.15& 22.79 $\pm$ 0.22& 3.37 $\pm$ 0.22 & 0.36 $\pm$ 0.27\\ 6 & 0.3076& 13.1630& 17.48 $\pm$ 0.00& 20.19 $\pm$ 0.03& 20.20 $\pm$ 0.05& 2.72 $\pm$ 0.05 & 0.01 $\pm$ 0.06\\ 7 & 0.3176& 13.0290& 18.43 $\pm$ 0.00& 22.85 $\pm$ 0.11& 21.70 $\pm$ 0.12& 3.27 $\pm$ 0.12 & -1.15 $\pm$ 0.16\\ 8 & 0.3188& 13.0724& 20.39 $\pm$ 0.01& 22.49 $\pm$ 0.13& 21.18 $\pm$ 0.11& 0.79 $\pm$ 0.11 & -1.31 $\pm$ 0.17\\ 9 & 0.3429& 13.1208& 20.35 $\pm$ 0.01& 22.24 $\pm$ 0.11& 22.10 $\pm$ 0.16& 1.75 $\pm$ 0.16 & -0.15 $\pm$ 0.19\\ 10& 0.3535& 13.0704& 19.75 $\pm$ 0.01& 21.92 $\pm$ 0.09& 21.66 $\pm$ 0.11& 1.91 $\pm$ 0.11 & -0.26 $\pm$ 0.14\\ 11& 0.3627& 13.0532& 19.42 $\pm$ 0.01& 21.63 $\pm$ 0.07& 21.62 $\pm$ 0.10& 2.20 $\pm$ 0.10 & -0.01 $\pm$ 0.12\\ 12& 0.3704& 13.1375& 18.88 $\pm$ 0.00& 22.15 $\pm$ 0.09& 21.98 $\pm$ 0.14& 3.10 $\pm$ 0.14 & -0.18 $\pm$ 0.17\\ 13$^{\rm b}$& 0.3735& 13.0986& 17.69 $\pm$ 0.00& 19.63 $\pm$ 0.02& 19.59 $\pm$ 0.04& 1.90 $\pm$ 0.04 & -0.04 $\pm$ 0.04\\ 13A& 0.3738& 13.0987& 17.83 $\pm$ 0.01& 19.78 $\pm$ 0.05& 19.81 $\pm$ 0.17& 1.98 $\pm$ 0.17 & 0.03 $\pm$ 0.18\\ 13B& 0.3719& 13.0984& 17.77 $\pm$ 0.01& 19.62 $\pm$ 0.04& 19.83 $\pm$ 0.04& 2.06 $\pm$ 0.04 & 0.22 $\pm$ 0.06\\ 14& 0.3773& 13.1041& 20.42 $\pm$ 0.01& 21.63 $\pm$ 0.08& 21.10 $\pm$ 0.09& 0.69 $\pm$ 0.09 & -0.53 $\pm$ 0.12\\ 15& 0.3796& 13.0950& 19.52 $\pm$ 0.01& 22.04 $\pm$ 0.12& 21.79 $\pm$ 0.14& 2.27 $\pm$ 0.14 & -0.25 $\pm$ 0.19\\ 16& 0.3816& 13.0817& 18.89 $\pm$ 0.00& 21.55 $\pm$ 0.06& 21.30 $\pm$ 0.11& 2.41 $\pm$ 0.11 & -0.24 $\pm$ 0.12\\ \enddata \tablenotetext{a}{Magnitudes (AB) in all bands were obtained with SExtractor Mag$_{-}$auto. Galactic extinction corrections were done using Schlegel et al. (1998) for R and Seibert et al. (2005) for FUV and NUV.} \tablenotetext{b}{Object 13 was separated into two objects, A and B, using IRAF polyphot task.} \end{deluxetable} \begin{deluxetable}{cccccccccc} \tabletypesize{\scriptsize} \tablecaption{Derived Properties} \tablewidth{0pt} \tablehead{ \colhead{ID} & \colhead{L$_{\rm FUV}$ (erg/s)$^{\rm a}$} & \colhead{L$_{\rm NUV}$ (erg/s)} & \colhead{age$^{\rm b}$} & \colhead{SFR$_{\rm FUV}$ $^{\rm c}$} & \colhead{SFR$_{\rm NUV}$} & \colhead{log(M$_{\rm *}$)$^{\rm d}$} & \colhead{log(M$_{\rm HI}$)$^{\rm e}$} & \colhead{Sep$^{\rm f}$} & \colhead{log I$_{1530}$$^{\rm g}$} } \startdata 1 & 5.90E+040 & 5.70E+040 & 3.9 & 0.005& 0.007 & 5.0 & & 131.2 & 6.02 \\ 2 & 2.27E+040 & 2.16E+040 & 3.8 & 0.002& 0.003 & 4.6 & & 102.2 & 5.91 \\ 3 & 2.91E+040 & 2.17E+040 & $<$1 & 0.002& 0.003 & 4.7 & 9.6 & 137.2 & 5.90 \\ 4 & 3.25E+040 & 2.98E+040 & 3.3 & 0.003& 0.004 & 4.7 & 10.4 & 60.7 & 5.35 \\ 5 & 6.79E+039 & 1.02E+040 & 194.1 & 0.001& 0.001 & 6.5 & 9.9$\dagger$ & 61.2 & 4.99 \\ 6 & 7.41E+040 & 8.01E+040 & 26.9 & 0.006& 0.010 & 6.3 & 9.2$\dagger$ & 74.7 & 6.31 \\ 7 & 1.84E+040 & 6.90E+039 & $<$1 & 0.001& 0.001 & 4.5 & & 108.5 & 5.51 \\ 8 & 2.98E+040 & 9.62E+039 & $<$1 & 0.002& 0.001 & 4.7 & 10.0 & 53.9 & 5.31 \\ 9 & 1.28E+040 & 1.21E+040 & 3.5 & 0.001& 0.001 & 4.6 & 9.8$\dagger$ & 17.5 & 5.80 \\ 10 & 1.92E+040 & 1.63E+040 & 2.9 & 0.002& 0.002 & 4.4 & & 58.8 & 5.97 \\ 11 & 2.00E+040 & 2.13E+040 & 21.0 & 0.002& 0.003 & 5.6 & & 83.7 & 5.85 \\ 12 & 1.43E+040 & 1.31E+040 & 3.2 & 0.001& 0.002 & 4.3 & & 58.0 & 5.95 \\ 13 & 1.29E+041 & 1.35E+041 & 13.7 & 0.010& 0.016 & 6.2 & 9.5$\dagger$$\dagger$ & 53.3 & 6.67 \\ 13 A& 1.03E+041 & 1.36E+041 & 118.7 & 0.008& 0.017 & & & & 6.51 \\ 13 B& 1.06E+041 & 1.17E+041 & 33.8 & 0.008& 0.014 & & & & 6.50 \\ 14 & 3.22E+040 & 2.13E+040 & $<$1 & 0.003& 0.003 & 4.8 & 9.5$\dagger$$\dagger$ & 56.3 & 6.30 \\ 15 & 1.70E+040 & 1.46E+040 & 2.9 & 0.001& 0.002 & 4.3 & 9.5$\dagger$$\dagger$ & 62.1 & 5.07 \\ 16 & 2.67E+040 & 2.29E+040 & 2.9 & 0.002& 0.003 & 4.5 & & 72.0 & 5.26 \\ \enddata \tablenotetext{a}{FUV and NUV luminosities are in erg/s, divide by the FUV and NUV bandwidths (269\AA\ and 616\AA, respectively) to obtain in L erg/s/\AA.} \tablenotetext{b}{Age (Myr) from FUV-NUV using Thilker et al. (2007) assuming a Milky Way internal extinction (E(B-V)=0.2).} \tablenotetext{c}{SFR (M$_{\odot}$/yr) from Iglesias-P\'aramo et al. (2006) using FUV and NUV without correcting for internal extinction.} \tablenotetext{d}{Stellar mass (M$_{\odot}$) obtained from Starburst99 using ages (column 4) and L$_{\rm 1530}$ (erg/s/\AA).} \tablenotetext{e}{MHI (M$_{\odot}$) was calculated using 2.36 $\times$ 10$^{5}$ F$_{\rm HI}$ D$^{2}$, where D is in Mpc and F$_{\rm HI}$ in Jy Km/s. F$_{\rm HI}$ was measured within one beam size 61.0$''$ $\times$ 55.23$''$ and reflects the HI mass in the vicinity of each object.} \tablenotetext{f}{Distance (kpc) between each object and HCG100a.} \tablenotetext{g}{FUV surface brightness (L$_{\odot}$ kpc$^{-2}$) is defined following Hoopes at al. (2007).} \tablenotetext{\dagger}{HI masses of Objects \#5, \#6, and \#9 should be taken with caution due to contamination by large galaxies in the vicinity.} \tablenotetext{\dagger\dagger}{HI masses of Objects \#13, \#14, and \#15 are not individual masses of each object but the total mass within one beam centered in that region. Contamination from H100b mass is also possible.} \end{deluxetable} | We have analyzed the UV and optical light in combination with the H{\sc i} gas within the compact group of galaxies HCG100 and identified 16 star-forming regions in the intragroup region. The young age ($<200$~Myr) of these objects and the proximity to the tidal tail connects the OB stars formation time scale ($\sim$10$^{8}$ yr) with the dynamic time scale of the tidal features. Moreover, the H{\sc i} clouds to which many of them are associated have column densities about one order of magnitude lower than the N(HI)$\sim$10$^{21}$ cm$^{-2}$ thought to be required for triggering star formation. So, in these cases, we have a strong indication that the H{\sc i} clouds must have suffered recent collisions which could have then triggered the star formation process. Based on their ages, stellar masses and H{\sc i} masses in their vicinities, we conclude that some of these objects are tidal dwarf galaxies with ongoing star formation and some are intergalactic HII regions or conglomeration of stellar clusters. Although we have an estimate of the amount of neutral gas in these objects, the main ingredient in the star-formation process, the molecular gas, is unknown for these objects. From CO observations of tidal dwarf galaxies (TDGs), Braine et al. (2001) provided strong evidence that TDGs are self-gravitating entities with large amounts of atomic gas which will transform into molecular gas and subsequently form stars. CO observations together with spectroscopy of these objects will be of great value in understanding their nature. For instance, we cannot exclude the possibility that some of these objects are not associated with HCG100, but at least those falling within HI density peaks are likely to be associated with the group. Moreover, we will be able to estimate metallicity of these newly formed regions and evaluate the enrichment level of the processed gas in the intragroup medium. \begin{deluxetable}{lcccc} \tablecaption{Comparison of Galaxy Properties} \tablewidth{0pt} \tablehead{ \colhead{Parameter$^{\rm a}$} & \colhead{Large UVLGs} & \colhead{Compact UVLGs} & \colhead{LBGs} & \colhead{HCG100 IGs} } \startdata log L$_{1530}$ (L$_{\odot}$) & 10.3 to 11.2 & 10.3 to 11.0 & 10.3 to 10.9 & 8 to 9 \\ log I$_{1530}$ (L$_{\odot}$ kpc$^{-2}$) $^{\rm b}$& 6.0 to 8.0 & 8.0 to 10.3 & 9.0 to 10.3 & 5.6 to 8.3 \\ log SFR (M$_{\odot}$ yr$^{-1}$) $^{\rm c}$ & 0 to 1.5 & 0.2 to 2.0 & 0.5 to 2.0 & -2 to -0.4 \\ FUV-R $^{\rm d}$ & 1.0 to 3.5 & 0.2 to 2.8 & 0.2 to 1.7 & 0.7 to 3.4 $^{\rm f}$ \\ \enddata \tablenotetext{a}{Parameters for UVLGS and LBGs are from Hoopes et al. (2007). The lowest value for HCG100 intragroup objects was calculated using average values A$_{1500}$=1.6 and the highest value using A$_{1500}$=4.1 from Calzetti (2001).} \tablenotetext{b}{I$_{1530}$ is the FUV surface brightness as described in Hoopes et al. (2007), except for HCG100 intragroup objects where we used R as the FUV petrosian radius.} \tablenotetext{c}{SFR for the HCG intragroup objects using FUV equation from Iglesias-P\'aramo et al. (2006).} \tablenotetext{d}{r magnitudes are from SDSS and corrected for extinction, except for HCG100 which is the same as in Table 1 (Cousin R magnitude) and is not corrected for extinction.} \tablenotetext{e}{FUV-R shown is not corrected for internal extinction. Using Calzetti (2001) A$_{1500}$=1.6 and A$_{r}$=0.48 FUV-R=-0.3 to 2.2.} \end{deluxetable} \clearpage \begin{figure} \plotone{f1.eps} \caption{GALEX FUV image 10$'$$\times$10$'$ with H{\sc i} contours. Four HCG100 members are marked as `a, b, c and d'. Intragroup objects are circled (radius=8$''$) and numbered. VLA NH{\sc i} contours are 0.6, 1.2, 2.1, 3.6, 4.4, 5.1, 5.9, 6.6, 7.4 $\times$10$^{20}$ cm$^{-2}$, beam size (61.0$''$$\times$55.23$''$) is shown on left corner. North is to the top and East to the left. \label{fuvhi_10arcmin}} \end{figure} \begin{figure} \plotone{f2.eps} \caption{GALEX NUV image 10$'$$\times$10$'$ with H{\sc i} contours. Intragroup objects are circled (radius=8$''$) and numbered. VLA NH{\sc i} contours are 0.6, 1.2, 2.1, 3.6, 4.4, 5.1, 5.9, 6.6, 7.4 $\times$10$^{20}$ cm$^{-2}$. Beam size (61.0$''$$\times$55.23$''$) is shown in Fig.\ref{fuvhi_10arcmin}. North is to the top and East to the left. \label{nuvhi_10arcmin}} \end{figure} \begin{figure} \plotone{f3.eps} \caption{GALEX FUV-NUV versus age from Thilker et al. models (2007) are shown as solid line (no extinction correction) and dotted line (internal extinction of the Milky Way E(B-V)=0.2). 4 objects are not included, they have ages $<$ 1~Myr old. \label{colorsfuvnuvthilker}} \end{figure} \begin{figure} \plotone{f4.eps} \caption{GALEX FUV-NUV versus FUV-R of the intragroup objects. Squares and triangles are FUV--R colors assuming an error of 0.2 magnitudes in the R-magnitude. Models from Thilker et al. 2007 are shown as solid line (no extinction correction) and dotted line (internal extinction of the Milky Way E(B-V)=0.2). Crosses mark ages 10, 50, 100, 200, and 500 Myr, from the bottom left to the top right. \label{colorsfuvnuvfuvr}} \end{figure} \begin{figure} \plotone{f5.eps} \caption{R images of all intragroup objects, circle has radius=8$''$, each cutout is 1$'$$\times$1$'$. NH{\sc i} contours (white) are the same as in Fig.\ref{fuvhi_10arcmin}. \label{all_obj}} \end{figure} \clearpage | 7 | 10 | 0710.3748 |
0710 | 0710.4397_arXiv.txt | We describe a finite-volume method for solving the Poisson equation on oct-tree adaptive meshes using direct solvers for individual mesh blocks. The method is a modified version of the method presented by Huang and Greengard (2000), which works with finite-difference meshes and does not allow for shared boundaries between refined patches. Our algorithm is implemented within the FLASH code framework and makes use of the PARAMESH library, permitting efficient use of parallel computers. We describe the algorithm and present test results that demonstrate its accuracy. | \label{Sec:intro} Astrophysical simulations commonly need to solve the Poisson equation, \begin{equation} \label{Eqn:Poisson} \nabla^2 \phi({\bf x}) = 4 \pi G \rho({\bf x})\ , \end{equation} for the gravitational potential $\phi({\bf x})$ given a density distribution $\rho({\bf x})$. Similar equations also arise in other contexts, such as incompressible flow problems and divergence-cleaning methods for magnetohydrodynamics. Self-gravitating problems offer special challenges because they frequently develop structure spanning large spatial dynamic ranges. The problem of spatial dynamic range is particularly acute for grid-based schemes for solving the Euler equations of hydrodynamics. Within the context of grid-based methods for solving the Poisson equation, several approaches to the problem of spatial dynamic range have arisen. The simplest approach is to use Fourier transforms, multigrid methods, or sparse iterative solvers on uniform Eulerian grids. The maximum dynamic range is then limited by the available memory. Recently Trac and Pen (2006) have demonstrated an out-of-core uniform-grid Poisson solver that exceeds this limit by making use of disk space; the largest published calculations with this solver have used $4000^3$ zones. However, storage resource consumption still increases with the third power of the resolution, putting grids with $10^4$ zones on a side or larger out of reach for now. If high resolution is not needed everywhere in the domain, as is frequently the case in cosmological structure formation simulations, it is also possible to employ nonuniform Eulerian or Lagrangian grids. Examples include COSMOS (Ricker, Dodelson, \& Lamb 2000), which uses a nonuniform multigrid solver, and MMH (Pen 1998), which uses a deformable mesh. These methods work best when the region to be resolved is known beforehand, although fully Lagrangian codes like Pen's can follow the development of structures and adjust zone spacing appropriately. Nonuniform grids, however, introduce complicated position-dependent stencils and generally cannot be used with fast transform-based solvers. In addition, coupled numerical hydrodynamics methods generally place constraints on the allowed mesh anisotropy and nonuniformity, since numerical dissipation increases with zone spacing. The greatest spatial dynamic ranges in grid-based astrophysical simulations have been achieved using adaptive mesh refinement (AMR) techniques. Modern AMR techniques for solving hyperbolic systems of equations were first developed by Berger and Oliger (1984) and Berger and Colella (1989). In the Berger and Colella formulation, AMR involves the construction of a hierarchical set of mesh ``patches'' with decreasing zone spacing. The coarsest mesh covers the entire computational domain, while more highly refined meshes cover only a portion. Generally refined meshes are taken to be nested; that is, each refined mesh lies completely within its coarser parent mesh. Examples of astrophysical codes employing patch-based AMR meshes include the code of Truelove et al.\ (1998), AMRA (Plewa \& M\"uller 2001), RIEMANN (Balsara 2001), Enzo (O'Shea et al.\ 2004), and CHARM (Miniati \& Colella 2007). To date self-gravitating AMR calculations have achieved effective spatial resolutions greater than $10^{15}$. A considerable simplification of the Berger and Colella method was introduced by Quirk (1991) and de~Zeeuw and Powell (1993). Known as ``oct-tree'' AMR, this method requires that each refined patch contain the same number of zones, that each refinement level have zones a factor of two smaller in each dimension than the next coarser level, and that each refined patch be no more than one level removed from its immediate neighbor. Mesh data can then be stored in an oct-tree data structure, allowing for extremely efficient parallel implementations, even on high-latency systems (Warren \& Salmon 1993). Also, because each mesh patch (often called a ``block'' in this context) contains the same number of zones, it is possible to achieve high levels of cache re-use when iterating over zones. Unless each block contains a very small number of zones, this efficiency comes with the price that refined blocks often must cover more of the volume than they would in a patch-based method. The primary astrophysical simulation code employing oct-tree AMR is FLASH (Fryxell et al.\ 2000), which uses the PARAMESH library (MacNeice et al.\ 2000) to handle its AMR mesh. (The ART (Kravtsov, Klypin, \& Khoklov 1997), MLAPM (Knebe, Green, \& Binney 2001), and RAMSES (Teyssier 2002) codes also use tree structures to manage AMR meshes, but in these codes the refined blocks consist of a single zone each, and the base mesh generally contains a large number of zones, so these codes do not employ oct-trees. A block-based AMR approach that allows for ``incomplete families'' has also recently been implemented within the VAC code (van der Holst \& Keppens 2007); the oct-trees discussed in the current paper require complete ``families'' of child blocks.) By default PARAMESH uses blocks containing $8^3$ interior zones as a compromise between adaptive flexibility and memory efficiency, but any size larger than the differencing stencil and small enough to fit in the memory attached to a single processor can be employed. Poisson solvers on oct-tree meshes generally employ some type of multigrid or sparse linear solver iteration scheme. Transform methods cannot be employed directly because of the varying mesh resolution and non-tensor-product character of the composite mesh. For example, Matsumoto and Hanawa (2003) describe a relaxation-based method for solving the Poisson equation on nested grids. Within the FLASH framework, we have employed the Martin and Cartwright (1996) multigrid algorithm for several years. However, the speed and scalability of this algorithm have been limited because of the need to apply multiple relaxation iterations on each level, together with the communication of block boundary data that such iterations require. Because the cost of this algorithm dominates the cost of most self-gravitating simulations with FLASH, we are motivated to develop more efficient methods that require less communication. In this paper we describe one such method, based on the direct multigrid algorithm of Huang and Greengard (2000, hereafter HG). This algorithm improves considerably upon relaxation-based multigrid solvers by allowing refined patches to be solved directly using ``black-box'' uniform-grid solvers. Unlike the direct method described by Couchman (1991), the HG algorithm properly minimizes the global residual by allowing information to flow back from fine meshes to coarse meshes. However, it is formulated on a finite-difference mesh in which refined patches are not permitted to touch. Here we describe a modified version of the HG algorithm suitable for finite-volume oct-tree AMR meshes. We have implemented this algorithm within the FLASH framework, and we present test results that demonstrate the solver's accuracy. The present paper should be regarded as a companion to Fryxell et al.\ (2000) and a methodological description for future FLASH-based simulation papers in the areas of cosmic structure formation, galaxy cluster physics, star cluster formation, and binary star evolution, among others. The paper is organized as follows. In \S~\ref{Sec:algorithm} we give a precise description of the algorithm. In \S~\ref{Sec:tests} we present the results of test problems run with the new solver. We conclude in \S~\ref{Sec:conclusions} with some remarks on performance. All calculations described in this paper were performed using version 2.4 of FLASH. | \label{Sec:conclusions} We have detailed the modifications needed in order to use the Huang \& Greengard (2000) algorithm for Poisson's equation on oct-tree AMR meshes. This method allows us to use a local direct Poisson solver with Dirichlet boundary conditions on each block, yet it correctly minimizes the residual across the composite AMR mesh. The HG algorithm must be modified to take into account the fact that mesh quantities represent zone-averaged values and the fact that block boundaries can coincide on an oct-tree mesh. Because adjacent coarse and fine blocks do not share points in common for oct-tree meshes, additional interpolation (as compared to HG) is needed to set boundary values. The resulting errors at block corners reduce convergence to first order, but a fixed small number of boundary-zone relaxation steps restores the desired second-order convergence rate for block corners in uniformly refined regions. A higher-order scheme may be able to eliminate the need for such relaxation. Our test results show that, while jumps in refinement degrade convergence somewhat in comparison with solving on a uniform mesh, the effect is manageable because oct-tree meshes are usually fully refined up to some level. Thus first-order convergence takes over only once the error has been significantly reduced at second order on fully refined levels. The parallel scaling of this solver, as implemented using the Message Passing Interface (MPI) in the FLASH code, is comparable to or slightly better than that of the relaxation multigrid solver distributed with FLASH 2.$x$. With constant total work and increasing processor count, parallel efficiency is close to 100\% up to 8 -- 16 times the smallest number of processors on which a run can fit. When gasdynamics is included, the Poisson solver requires $\sim$50\% of the execution time; for particle-only simulations the amount is $\sim $70 -- 80\%. In comparison with the older solver, a factor of 2 -- 5 improvement in performance is often seen. The solver described here will be made available to the public as part of FLASH~3.0.\footnote{FLASH is freely available at http://flash.uchicago.edu/.} The primary scaling bottleneck for this and other multigrid algorithms is the fact that on the coarsest level there are too few zones to distribute among all of the processors. Since each V-cycle descends to the coarsest level, processors controlling blocks on finer levels must wait until the coarsest level is finished before proceeding. To counteract this work starvation, we are investigating the use of a uniform-grid parallel FFT solver (T.~Theuns, private communication) to handle the coarsest level in a distributed fashion. Additional strategies may be necessary to make optimal use of the petaflop computing resources now beginning to become available. | 7 | 10 | 0710.4397 |
0710 | 0710.4674_arXiv.txt | We present a study of large-scale bars in the local Universe, based on a large sample of $\sim3692$ galaxies, with $-18.5\leq M_g < -22.0$ mag and redshift $0.01\leq z<0.03$, drawn from the Sloan Digitized Sky Survey. While most studies of bars in the local Universe have been based on relatively small samples that are dominated by bright galaxies of early to intermediate Hubble types with prominent bulges, the present sample is $\sim$~10 times larger, covers a larger volume, and includes many galaxies that are disk-dominated and of late Hubble types. Both color cuts and S\'ersic cuts yield a similar sample of $\sim2000$ disk galaxies. We characterize bars and disks by ellipse-fitting $r$-band images and applying quantitative criteria. After excluding highly inclined ($>60^{\circ}$) systems, we find the following results. (1)~The optical $r$-band fraction ($f_{\rm opt-r}$) of barred galaxies, when averaged over the whole sample, is $\sim48\%-52\%$. The bars have diameters $d$ of 2 to 24 kpc, with most ($\sim72\%$) having $d\sim$ 2 to 6 kpc. (2)~When galaxies are separated according to half light radius ($r_{\rm e}$), or normalized $r_{\rm e}$/$R_{\rm 24}$, which is a measure of the bulge-to-disk ($B/D$) ratio, a remarkable result is seen: $f_{\rm opt-r}$ rises sharply, from $\sim$~40\% in galaxies that have small $r_{\rm e}$/$R_{\rm 24}$ and visually appear to host prominent bulges, to $\sim$~70\% for galaxies that have large $r_{\rm e}$/$R_{\rm 24}$ and appear disk-dominated. Visual classification, performed for $\sim900$ galaxies, confirms our result that disk-dominated galaxies with no bulge or a very low $B/D$ display a significantly higher optical bar fraction ($>70$\% vs 40\%) than galaxies with prominent bulges. It also shows that barred galaxies host a larger fraction (31\% vs 5\%) of quasi-bulgeless disk-dominated galaxies than do unbarred galaxies. (3)~$f_{\rm opt-r}$ rises for galaxies with bluer colors (by $\sim30\%$) and lower masses (by $\sim15\%-20\%$). (4) The significant rise in the optical bar fraction toward late-type galaxies is discussed in terms of their higher gas mass fraction, higher dark matter fraction, and lower bulge-to-disk ratio. (5) While hierarchical $\Lambda$CDM models of galaxy evolution models fail to produce galaxies without classical bulges, our study finds that $\sim20\%$ of disk galaxies appear to be ``quasi-bulgeless''. (6)~Our study of bars at $z\sim0$ in the optical $r$-band provides the $z\sim0$ comparison point for $HST$ ACS surveys (e.g., GEMS, GOODS, COSMOS) that measure the rest-frame optical bar fraction in bright galaxies out to $z\sim1$. After applying the same cutoffs in magnitude, bar ellipticity ($e_{\rm bar}\ge0.4$), and bar size ($a_{\rm bar}\ge1.5$ kpc), which are applied in $z\sim0.2-1.0$ studies in order to trace strong bars with adequate spatial resolution in bright disks, we obtain an optical $r$-band bar fraction of $34\%$. This is comparable to the value reported at $z\sim0.2-1.0$, implying that the optical fraction of strong bars does not suffer a dramatic order of magnitude decline in bright galaxies out to $z\sim1$. | The majority ($\sim$~60\%) of bright disk galaxies are barred, when observed in the near-infrared \citep[][hereafter MJ07]{kna99,esk00,lau04,men07,mar07} and a significant fraction of these ($\sim$~45\%) also appear barred in the optical \citep[][MJ07]{esk00}. Earlier studies suggested a striking or order of magnitude decline in the optical fraction of bars out to $z\sim1$ \citep{abr99,vdb00}, but subsequent studies have ruled out an order of magnitude decline and find that the optical fraction of strong bars remains fairly constant or show a moderate decline of a factor of $\sim2$ \citep[][Sheth et al. 2003,2004,2007; see $\S$ \ref{discu}]{jog04,elm04a,zhe05}. Bars are believed to be very important with regard to the dynamical and secular evolution of disk galaxies, particularly in redistributing the angular momentum of the baryonic and dark matter components of disk galaxies \citep{com81,wei85,com90,deb00}. The interaction between the bar and the disk material can lead to the inflow of gas from the outer disk to the central parts, which can trigger starbursts \citep{elm94,kna95,reg04,jog05,sht05} and might contribute to the formation of disky bulges \citep[or 'pseudobulges',][]{kor93,kor04,ath05a,jog05,deb06}. Additional evidence for secular evolution is provided by box- or peanut-shaped bulges in inclined galaxies. These features are commonly attributed to the orbital structure, resonances, and vertical instabilities in a barred potential \citep{com90,kui95,bur99,bur05}. From a theoretical perspective, it is possible to model some aspects of the evolution of disks and bars, and their interactions (e.g., the corresponding simulations are able to reproduce certain broad features of barred disks). However, it remains unclear why a specific galaxy has a bar, but a seemingly similar galaxy is unbarred; or why some barred galaxies have a classical bulge, whereas others harbor a disky bulge, etc. This might indicate that specific properties of the disks or the particular processes involved in their formation have a strong impact on their ability to form a bar. In order to investigate, how disk and bar formation are related, it is not only important to determine the fraction of disk galaxies that are barred, but also to relate bar and disk properties. There are different methods to find and characterize bars. The Third Reference Catalog of Bright Galaxies \citep[][hereafter RC3]{dev91} uses three bar strength families (SA, SAB, and SB) to characterize bars based on a visual inspection of blue light images. Using this classification \cite{ode96} showed that the optical fraction of strong bars in disk galaxies rises from Sc galaxies towards later-types. More quantitative measures, such as the gravitational torque method \citep{blo02,lau02,but05}, or Fourier dissection \citep{but06,lau06}, were also used, not only to find bars, but to quantitatively determine bar strengths and bar lengths. Similarly, the method of fitting ellipses to galaxy isophotes provides a tool to characterize the length and shape of bars \citep{fri96,jog99,kna00,sht00,lai02,why02,jog02a,jog02b,sht03,elm04a,ree07,mar07,men07,sht07}. These efforts were able to shed light on the fraction, shapes, and structures of bars in local disk galaxies of early to intermediate Hubble types. First attempts were made to relate the presence of a bar or its structural properties to other galaxy characteristics. However, there were three important limitations. Firstly, samples used in earlier studies were small ($\sim100$ to 200 objects) and mostly composed of bright galaxies of early to intermediate Hubble types (Sa to Sc), with fairly prominent bulges. One could barely get decent number statistics for bars in early-type disk galaxies, while the bins of disk-dominated late Hubble types were dominated by Poisson noise (e.g., see Figure 16 in MJ07). Secondly, with such small samples, it was difficult to bin galaxies in terms of the galaxy host properties. Thirdly, earlier samples were drawn from a very small volume, and could be highly impacted by cosmic variance. In the present study, we use a sample of $\sim2000$ galaxies, at $z=0.01-0.03$ with $M_r \sim -18.5$ to $-22.0$ mag. The first advantage of this study is that it provides a factor of 10 improvement in number statistics and reduces the effect of cosmic variance by selecting galaxies drawn from a larger volume. Secondly, with $\sim2000$ galaxies, we can for the first time have $100-200$ galaxies per bin, while binning galaxies in terms of host galaxy parameters, such as luminosity, measures of bulge-to-disk ($B/D$) ratios, size, colors, surface brightness, etc. This allows us to conduct a comprehensive study of barred and unbarred galaxies as a function of host galaxy properties. Thirdly, our sample has a large number of galaxies, which are relatively faint ($M_g>-19.5$ mag) or/and appear disk-dominated, characteristic of late Hubble types. This allows us to shed light on what happens to bars at the fainter end of the luminosity function and in the regime of disk-dominated galaxies. A fourth goal of our study is to provide a reference baseline for bars at $z\sim0$ in the {\it rest-frame optical} for intermediate redshift $HST$ surveys using the Advanced Camera for Surveys (ACS), such as the Tadpole field \citep{tra03}, the Galaxy Evolution from Morphologies and SEDs \citep[GEMS,][]{rix04}, the Great Observatories Origins Deep Survey \citep[GOODS,][]{gia04}, and COSMOS \citep{sco06}, which trace bars in the rest-frame optical band at $z\sim0.2-1.0$ (look-back times of 3--8 Gyr). We use SDSS to provide the reference point at $z=0$ in the $r$-band, complementing the one in the $B$-band of MJ07. Our $B$- and $r$-band results can be directly compared to $HST$ ACS optical studies of bars in bright disks at $z\sim0.2-1.0$ \citep{elm04a,jog04}. The validity of this comparison is reinforced by the fact that we use the same procedure of ellipse fits ($\S$ \ref{meth}) that were used by these studies. We also note that the reference $z=0$ point for bars in the near-infrared band \citep{men07} is not appropriate for comparison with the above $HST$ ACS surveys, which trace the rest-frame optical rather than the rest-frame near-infrared. The outline of the paper is as follows: In $\S$ \ref{samsel} we present our sample selection. The method used to find and characterize bars is explained in $\S$ \ref{meth}. In $\S$ \ref{resol} we discuss the detection limits. Our results and more detailed assessments of specific findings are presented in $\S$ \ref{resul}. We discuss our results in $\S$ \ref{discu} and summarize our conclusions in $\S$ \ref{sum}. Throughout the paper we assume a flat cosmology with $\Omega_M=1-\Omega_{\Lambda}=0.3$ and $H_0=70$ km~s$^{-1}$ Mpc$^{-1}$. | \label{sum} We have used the $r$-band images from the NYU-VAGC of a sample of 3692 galaxies with $-18.5\leq M_g < -22.0$ mag and redshift $0.01\leq z<0.03$ to find and characterize bars. While most studies of bars in the local Universe have been based on relatively small samples that are dominated by bright early type (Sa to Sc) galaxies with bulges, the present sample also includes many galaxies that are disk-dominated and of late Hubble types. Furthermore, the sample is $\sim$~10 times larger and samples a larger volume than earlier local samples We used a color cut in the color-magnitude diagram to select $\sim2000$ disk galaxies. We cross-check that S\'ersic cuts would yield a similar sample. We identify and characterize bars and disks using $r$-band images and a method based on ellipse fits and quantitative criteria. The typical seeing ($1\farcs4$ or 290 to 840 pc over $0.01\leq z<0.03$) is adequate for resolving large-scale bars, whose typical diameters are $\ge$ 2 kpc. Smaller nuclear bars are not the focus of this study. After the standard procedure of excluding highly inclined ($>60^{\circ}$) systems, we find the following results. \begin{enumerate} \item The average optical $r$-band bar fraction ($f_{\rm opt-r}$) in our sample, which primarily consists of late-type disk-dominated galaxies, is $\sim48\%-52\%$. The bars have diameters $d$ of 2 to 24 kpc, with most ($\sim72\%$) having $d\sim$ 2 to 6 kpc (Figure \ref{barp}a). The bar length is typically much smaller than $R_{24}$ (Figure \ref{r24}a) and most galaxies have a $a_{bar}/R_{24}$ in the range 0.2 to 0.4 (Figure \ref{r24}b). \item When galaxies are separated according to normalized $r_{\rm e}$/$R_{\rm 24}$, which is a measure of the bulge-to-disk ($B/D$) ratio, a remarkable result is seen: the optical $r$-band fraction rises sharply, from $\sim$~40\% in galaxies that have small $r_{\rm e}$/$R_{\rm 24}$ and visually appear bulge-dominated, to $\sim$~70\% for galaxies that have large $r_{\rm e}$/$R_{\rm 24}$. Visual classification of $\sim80\%$ of our sample (with $i<60^{\circ}$) confirms our result that {\it late-type disk-dominated galaxies with no bulge or a very low $B/D$ display a significantly higher optical bar fraction ($>70$\% vs 40\%) than galaxies with prominent bulges.} It also shows that barred galaxies host a larger fraction (31\% vs 5\%) of quasi-bulgeless disk-dominated galaxies than do unbarred galaxies. The bar ellipticities or strengths are on average higher in faint disk-dominated galaxies than in bulge-dominated galaxies (Figure \ref{visp2}d). \item Similar trends in the optical bar fraction are found using the central surface brightness and color. Bluer galaxies have higher bar fractions ($\sim58\%$ at $g-r=0.3$) than the redder objects ($\sim32\%$ at $g-r=0.65$) (Figure \ref{barf}d). The optical $r$-band fraction also shows a slight rise for galaxies with fainter luminosities (Figure \ref{barf}c) and lower masses (Figure \ref{massp}). This is expected from (2), given that late-type galaxies are fainter, bluer, and less massive. \item The significant rise in the optical bar fraction toward disk-dominated galaxies is discussed in terms of their higher gas mass fraction, higher dark matter fraction, and lower bulge-to-disk ratio. \item While many hierarchical $\Lambda$CDM models of galaxy evolution models fail to produce galaxies without classical bulges, our study finds that in the range $-18.5\leq M_g < -22.0$ mag and redshift $0.01\leq z<0.03$, $\sim20\%$ of the 1144 moderately inclined disk galaxies appear to be ``quasi-bulgeless'', without a classical bulge. \item Our study of bars at $z\sim0$ in the optical $r$ band provides a reference $z\sim0$ baseline for intermediate redshift $HST$ ACS surveys that trace bars in {\it bright} disks in the rest-frame optical bands ($BVRI$) out to $z\sim1$. By applying the same cutoffs in magnitude, bar ellipticity ($e_{\rm bar} \geq0.4$), and bar size ($a_{\rm bar} \geq1.5$ kpc), which are applied in $z\sim0.2-1.0$ studies in order to trace strong bars with adequate spatial resolution in bright disks, we obtain an optical $r$-band fraction for strong bars of $34\%$. This is comparable to the values of $\sim30\%$ at $z\sim0.2-1.0$, $\sim$~36\%~$\pm$~6\% at $z\sim0.2-0.7$, and $\sim$~24\%~$\pm$~4\% at $z\sim0.7-1.0$. Our result implies that the optical fraction of strong bars in bright galaxies does not suffer any dramatic order of magnitude decline out to $z\sim$~1. \end{enumerate} | 7 | 10 | 0710.4674 |
0710 | 0710.4442_arXiv.txt | We construct for the first time, the sequences of stable neutron star (NS) models capable of explaining simultaneously, the glitch healing parameters, $Q$, of both the pulsars, the Crab ($Q \geq 0.7$) and the Vela ($Q \leq 0.2$), on the basis of starquake mechanism of glitch generation, whereas the conventional NS models cannot give such consistent explanation. Furthermore, our models also yield an upper bound on NS masses similar to those obtained in the literature for a variety of modern equations of state (EOSs) compatible with causality and dynamical stability. If the lower limit of the observational constraint of (i) $Q \geq 0.7$ for the Crab pulsar and (ii) the recent value of the moment of inertia for the Crab pulsar (evaluated on the basis of time-dependent acceleration model of the Crab Nebula) , $I_{\rm Crab,45} \geq 1.93$ (where $I_{45}=I/10^{45}\,{\rm g.cm}^2$), both are imposed together on our models, the models yield the value of matching density, $E_b = 9.584 \times 10^{14}{\rm\,g\,cm}^{-3}$ at the core-envelope boundary. This value of matching density yields a model-independent upper bound on neutron star masses, $M_{\rm max} \leq 2.22 M_\odot$, and the strong lower bounds on surface redshift $z_R \simeq 0.6232$ and mass $M \simeq 2.11 M_\odot$ for the Crab ($Q \simeq 0.7$) and the strong upper bound on surface redshift $z_R \simeq 0.2016 $, mass $M \simeq 0.982 M_\odot$ and the moment of inertia $I_{\rm Vela,45} \simeq 0.587$ for the Vela ($Q \simeq 0.2$) pulsar. However, for the observational constraint of the `central' weighted mean value $Q \approx 0.72$, and $I_{\rm Crab,45} > 1.93$, for the Crab pulsar, the minimum surface redshift and mass of the Crab pulsar are slightly increased to the values $z_R \simeq 0.655$ and $M \simeq 2.149 M_\odot$ respectively, whereas corresponding to the `central' weighted mean value $Q \approx 0.12$ for the Vela pulsar, the maximum surface redshift, mass and the moment of inertia for the Vela pulsar are slightly decreased to the values $z_R \simeq 0.1645,\, M \simeq 0.828 M_\odot$ and $I_{\rm Vela,45} \simeq 0.459$ respectively. These results set an upper and lower bound on the energy of a gravitationally redshifted electron-positron annihilation line in the range of about 0.309 - 0.315 MeV from the Crab and in the range of about 0.425 - 0.439 MeV from the Vela pulsar. | The data on the glitch healing parameter, $Q$, provide the best tool for testing the starquake (Ruderman 1972; Alpar et al 1996) and Vortex unpinning (Alpar et al 1993) models of glitch generation in pulsars. Both of these mechanisms of glitch generation, in fact, consider NSs, in general, a two component structure: a superfluid interior core surrounded by a rigid crust (in the present study we shall use the term `envelope' which includes the solid crust and other interior portion of the star right up to the superfluid core). In the starquake model, $Q$ is defined as the fractional moment of inertia, i.e. the ratio of the moment of inertia of the superfluid core, $I_{\rm core}$, to the moment of inertia of the entire configuration, $I_{\rm total}$, as (Pines et al 1974) \begin{equation} Q = \frac{I_{\rm core}}{I_{\rm total}}. \end{equation} Recently, Crawford \& Demia\'{n}ski (2003) have collected the all measured values of the glitch healing parameter $Q$ for Crab and Vela pulsars available in the literature and found that for 21 measured values of $Q$ for Crab glitches, a weighted mean of the values yields $Q = 0.72 \pm 0.05$, and the range of $Q \geq 0.7$ encompasses the observed distribution for the Crab pulsar. In order to test the starquake model for the Crab pulsar, they have computed $Q$ (as given by Eq.(1)) values for seven representative EOSs of dense nuclear matter, covering a range of neutron star masses. Their study shows that the much larger values of $Q(\geq 0.7)$ for the Crab pulsar is fulfilled by all the six EOSs (out of seven considered in the study) corresponding to a `realistic' neutron star mass range $1.4\pm 0.2M_\odot$. By contrast, a weighted mean value of the 11 measurements for Vela yields a much smaller value of $Q(= 0.12 \pm 0.07)$ and the all estimates for Vela agree with the likely range of $Q \leq 0.2$. Thus, their results are found to be consistent with the starquake model predictions for the Crab pulsar. They have also concluded that the much smaller values of $Q \leq 0.2$ for the Vela pulsar are inconsistent with the starquake model predictions, since the implied Vela mass based upon their models corresponds to a value $ \leq 0.5M_\odot$ for $Q \leq 0.2$, which is too low as compared to the `realistic' NS mass range. Thus, in the literature, the starquake is considered as a viable mechanism for glitch generation in the Crab and the vortex unpinning, the another mechanism, is considered suitable for the Vela pulsar, since it can avoid some other problems associated with the starquake explanation of the Vela glitches (see, e.g. Crawford \& Demia\'{n}ski (2003); and references therein). However, it seems surprising that if the internal structure of NSs are described by the same two component conventional models (as mentioned above), different kinds of glitch mechanisms are required for the explanation of a glitch! Furthermore, it also follows from the above discussion that the main reason for not considering the starquake, the feasible mechanism for glitch generation in the Vela, lies in the fact that there exists none of the sequence of NS models in the literature which could explain simultaneously, on the basis of starquake model, both the extreme limiting cases of glitch healing parameter, $Q$, corresponding to the Vela ($Q \leq 0.2$) and the Crab ($Q \geq 0.7$) pulsars in the range $0 \leq Q \leq 1.0$ for the `realistic' NS mass values for both the pulsars. The present study, therefore, deals with the construction of such models \footnote{however, the other problems associated with the starquake explanation of the Vela glitches (see, e.g. Crawford \& Demia\'{n}ski (2003); and references therein) are not considered in the present paper. The future study in this regard may provide some explanation, provided the correlation between various parameters of the Crab and the Vela pulsar, obtained in the present study, can be utilized.} We assume that {\em all} the NSs belong to the same family of NS sequence which terminates at the {\em maximum} value of mass. Certainly, this {\em maxima} should correspond to an {\em upper bound} on NS masses. In order to construct such a sequence, we have to set the extreme causal EOS (in geometrized units), $dP/dE = 1$ (where $P$ is the pressure and $E$ the energy-density) to describe the core. Firstly, because various observational studies like - the gamma-ray burst data, X-ray burst data and the glitch data etc., and their explanation (see, e.g. Lindblom 1984; Cottam et al 2002; Datta \& Alpar 1993) favour the stiffest EOSs. The latest estimate of the moment of inertia for the Crab pulsar (based upon the `newest' observational data on the Crab nebula mass) rules out most of the existing EOSs of the dense nuclear matter, leaving only the stiffest ones (Bejger \& Haensel 2002; Haensel et al 2006). Secondly, because of the fact that the `real' EOS of the dense nuclear matter beyond the density range $\sim 10^{14}\,{\rm g\, cm}^{-3}$ are largely unknown due to the lack of knowledge of nuclear interactions (see, e.g. Dolan 1992; and references therein; Haensel et al 2006), and the various EOSs available in the literature (see, e.g. Arnett \& Bowers 1977) for NS matter represent only an extrapolation of the results far beyond this density range. Though, the status of the `real' EOS for NS matter is not certain, one could impose some well-known physical principle, independent of the EOS, such as the `causality condition' ($dP/dE = 1$) throughout the core of the star beyond a fiduciary density, $E_b$, at the core-envelope boundary to ascertain a definite upper bound on NS masses (see, e.g., Rhoades \& Ruffini 1974; Hartle 1978; Lindblom 1984; Friedman \& Ipser 1987; Kalogera \& Baym 1996). In this connection this is also to be pointed out here that the maximum mass for {\em any} EOS describing the core, beyond the density $E_b$, with a subluminal sound velocity turns out to be less than that of the upper bound obtained by using the extreme causal EOS (see, e.g., Haensel et al 2006). The envelope of our models (below the density $E_b$ at the core-envelope boundary) may be characterized by the well-known EOS of classical polytrope ${\rm d}$ln$P/{\rm d}$ln$\rho = \Gamma_1$ (where $\rho$ denotes the density of the rest-mass and $\Gamma_1$ is a constant known as the adiabatic index) for different values of the constant $\Gamma_1 = (4/3), (5/3)$ and 2 respectively. The reason for considering the polytropic EOS for the entire envelope lies in the fact that with this EOS, our models yield an upper bound on NS masses {\em independent} of the value of $\Gamma_1$, and this upper bound (for a fiduciary choice of $E_b$) is found fully consistent with those of the values cited in the literature (Kalogera \& Baym 1996; Friedman \& Ipser 1987). Thus, the choice of the said polytropic EOS for the entire envelope may be regarded entirely equivalent to the choice of the various EOSs like WFF (Wiringa, Fiks \& Fabrocini 1988), FPS (Lorenz, Ravenhall \& Pethick, 1993), NV (Negele \& Vautherin 1973), or BPS (Baym, Pethick \& Sutherland 1971) in an appropriate sequence below the density range $E_b$, adopted by various authors in the conventional models of NSs (see, e.g., Kalogera \& Baym 1996; Friedman \& Ipser 1987), so that the constant $\Gamma_1$ appearing in the polytropic EOS may be looked upon as an `average' $\Gamma_1$ for the density range below $E_b$, specified by the sequence of various EOSs in the conventional models of NSs. The choice of the constant $\Gamma_1 = 4/3, 5/3$ and 2 thus become obvious, since this choice can cover almost the entire range of density discussed in the literature for NS matter which is also applicable for the envelope region - the polytropic EOS with $\Gamma_1 = 4/3$ represents the EOS of extreme relativistic degenerate electrons and non-relativistic nuclei (Chandrasekhar 1935), $\Gamma_1 = (5/3)$ represents the well-known EOS of non-relativistic degenerate `neutron gas' (Oppenheimer \& Volkoff 1939), and $\Gamma_1 = 2$ represents the case of extreme relativistic baryons interacting through a vector meson field (Zeldovich 1962) (The value of $\Gamma_1 > 2$ is also possible for some EOS describing the NS matter, e.g., Malone, Johnson \& Bethe 1975; Clark, Heintzmann \& Grewing 1971, however, the results obtained in this paper remain unaffected for the choice of $\Gamma_1 > 2$), and the outcome of this study (in terms of explaining the glitch healing parameter for various pulsars and predicting the upper bound on the compactness of NSs (since the upper bound on mass is independent of the value of $\Gamma_1$)) would finally decide, among the chosen values, the `appropriate' value of $\Gamma_1$ for the NS envelope. The validity of assuming the extreme causal EOS in the core and a polytropic EOS in the envelope of the present models, in view of the various modern EOS of dense nuclear matter, is also discussed in the last section of the present paper. We have noted that in all conservative models of NSs, the choice of the core-envelope boundary, $r_b$ (corresponding to a density denoted by $E_b$), is somewhat {\em arbitrary} in the sense that there are no criteria available for the choice of a particular matching density, $E_b$, below which the EOS of the NS matter is assumed to be known and unique. One can freely choose somewhat lower values of $E_b$ (which will increase the core size) to obtain higher values of $Q$ (see, e.g. Shapiro \& Teukolsky 1983; Datta \& Alpar 1993). To avoid such a procedure, we choose the core-envelope boundary of our models on the basis of the `compatibility criterion' which asserts that for an assigned value of the ratio ($\sigma$) of central pressure, $P_0$, to central energy-density, $E_0$, the compactness parameter $u(\equiv M/R$; total mass to radius ratio in geometrized units) of any {\em regular} configuration should not exceed the compactness parameter $u_h$ of the homogeneous density sphere, in order to assure the compatibility with the hydrostatic equilibrium (Negi \& Durgapal 2001; Negi 2004a). This criterion is capable of constraining the core-envelope boundary of any physically realistic NS model. A combination of this criterion with those of the observational data on the glitch healing parameter and the recently estimated minimum value of the moment of inertia for the Crab pulsar (based on the newly estimated `central value' of the Crab nebula mass $M {\rm (nebula)} \simeq 4.6 M_\odot$ in the time-dependent acceleration model), $I_{\rm Crab,45} = 1.93$; where $I_{\rm Crab,45} = I_{\rm Crab}/10^{45}$ g\,cm$^2$ (Bejger \& Haensel 2003) can provide the desired NS models discussed above, since both the theory (criterion) and the observations (stated above) are being used to construct the NS models. | This study constructs the stable sequences of NS models terminate at the value of maximum mass, $M_{\rm max} \simeq 2.22 M_\odot$, independent of the EOSs of the envelope, for the matching density, $E_b = 9.584 \times 10^{14}$ g\, cm$^{-3}$, at the core-envelope boundary. This value of `matching density' is a consequence of the observational constraints $Q \simeq 0.7$ and $I_{\rm Crab,45} \simeq 1.93$ (associated with the Crab pulsar) imposed together on the $\Gamma_1 = 2$ envelope model and in this sense does not represent a fiduciary quantity. The upper bound of the surface redshift, $z_R \simeq 0.77$ (corresponding to a $u$ value $\simeq 0.34$), however, belongs to the model with a $\Gamma_1 = 2$ envelope which is consistent with the absolute upper bound on the surface redshift of NS models compatible with causality and pulsational stability (Negi 2004b). This special feature, together with some other remarkable ones, discussed in the present study underline the appropriateness of the $\Gamma_1 = 2$ envelope model. Since among the variety of modern EOSs discussed in the literature, the upper bound on NS mass compatible with causality and dynamical stability can reach a value up to $2.2M_\odot$ (in this category, the SLy (Douchin \& Haensel 2001) EOS yields a maximum mass of $2.05M_\odot$, whereas the BGN1 (Balberg \& Gal 1997) and the APR (Akmal et. al. 1998) EOSs yield the maximum masses of $2.18M_\odot$ and $2.21M_\odot$ respectively (see, e.g. Haensel et al 2006)). In view of this result, the model-independent maximum mass, $M_{\rm max} \simeq 2.22M_\odot$, obtained in this study may be regarded as good as those obtained on the basis of modern nuclear theory. In addition to this result, the appropriate sequences of stable NS models obtained in this study can explain the glitch healing parameter, $Q$, of any glitching pulsar, provided the weighted mean values of $Q$ lie in the range $0 < Q \leq 0.78$. This finding also reveals that if the starquake is considered to be a viable mechanism for glitch generation in all pulsars, then the envelope of `real' NSs may be well approximated by a polytropic EOS corresponding to a polytropic index, $n$, closer to 1. For the value of matching density, $E_b = 9.584 \times 10^{14}$ g\, cm$^{-3}$, the $\Gamma_1 = 2$ envelope model yields the minimum values of mass $M \simeq 2.11 M_\odot$ and surface redshift $z_R \simeq 0.6232$ for the Crab ($Q \simeq 0.7$) and the maximum values of mass $M \simeq 0.982 M_\odot$ and surface redshift $z_R \simeq 0.2016$ for the Vela pulsar ($Q \simeq 0.2$). The minimum mass and surface redshift for the Crab pulsar are slightly increased up to the values $M \simeq 2.149 M_\odot$ and $z_R \simeq 0.655$ respectively, if the `central' weighted mean value of $Q \approx 0.72$ and the moment of inertia $I_{\rm Crab,45} > 1.93 $ are also imposed on these models. However, for the `central' weighted mean value of $Q \simeq 0.12$ corresponding to the Vela pulsar, the maximum mass and surface redshift are somewhat decreased to the values $M \simeq 0.828 M_\odot$ and $z_R \simeq 0.1645$ respectively. This value of mass and surface redshift for the Vela pulsar can further decrease up to the values $M \simeq 0.685 M_\odot$ and $z_R \simeq 0.1312$ respectively, if the lower weighted mean value of $Q \simeq 0.05$ for the Vela pulsar is imposed. These results predict the upper and lower bounds on the energy of a gravitationally redshifted electron-positron annihilation line in the range of about 0.309 - 0.315 MeV from the Crab and in the range of about 0.425 - 0.439 MeV from the Vela pulsar respectively. For a comparison, if the observational constraint of the minimum value of $I_{\rm Crab,45} \simeq 3.04$ (the value of moment of inertia for the Crab pulsar obtained earlier by Bejger \& Haensel (2002), on the basis of the constant-acceleration model for the Crab nebula) together with $Q \simeq 0.7$ is imposed on the models studied in the present paper, the $\Gamma_1 = 2$ envelope model yields the value of matching density, $E_b = 7.0794 \times 10^{14}$ g\, cm$^{-3}$. This value of $E_b$ yields a model-independent upper bound on NS mass $M_{\rm max} \simeq 2.59 M_\odot$. This value of maximum mass, however, represents an `average' of the maximum NS masses in the range $2.2M_\odot \leq M_{\rm max} \leq 2.9M_\odot$ obtained by Kalogera \& Baym 1996 (and references therein) on the basis of other EOSs for NS matter, fitted to experimental nucleon-nucleon scattering data and the properties of light nuclei. For this lower value of matching density, the $\Gamma_1 = 2$ envelope models yield the minimum value of mass $M \simeq 2.455 M_\odot$ for the Crab ($Q \simeq 0.7$) and the maximum value of mass $M \simeq 1.142 M_\odot$ for the Vela pulsar ($Q \simeq 0.2$). The minimum mass for the Crab pulsar is slightly increased up to the value $M \simeq 2.5 M_\odot$, if the `central' weighted mean value of $Q \approx 0.72$ and the moment of inertia $I_{\rm Crab,45} > 3.04 $ are also imposed on these models. However, corresponding to the `central' weighted mean value of $Q \simeq 0.12$, the maximum mass of the Vela pulsar is somewhat decreased to the value $M \simeq 0.964 M_\odot$. This value of mass for the Vela pulsar can further decrease up to the value $M \simeq 0.796 M_\odot$, if the lower weighted mean value of $Q \simeq 0.05$ for the Vela pulsar is imposed. Furthermore, the study can also explain some special features associated with the extraordinary gamma-ray burst of 5 March 1979. | 7 | 10 | 0710.4442 |
0710 | 0710.4168_arXiv.txt | We have applied the torus fitting procedure described in Ng \& Romani (2004) to PWNe observations in the \emph{Chandra} data archive. This study provides quantitative measurement of the PWN geometry and we characterize the uncertainties in the fits, with statistical errors coming from the fit uncertainties and systematic errors estimated by varying the assumed fitting model. The symmetry axis $\Psi$ of the PWN are generally well determined, and highly model-independent. We often derive a robust value for the spin inclination $\zeta$. We briefly discuss the utility of these results in comparison with new radio and high energy pulse measurements. | One of the greatest success of the Chandra X-ray Observatory (CXO) is the discovery of equatorial tori and polar jet structures in many pulsar wind nebula (PWN) systems. It is now believe that these features are common among young neutron stars. In the \citet{ree74} picture, when the highly relativistic pulsar wind decelerates in the external medium, a termination shock is formed at a characteristic scale \[r_t =\left(\frac{\dot E}{4\pi c \eta P_{\rm ext}} \right )^{1/2} \ , \] where $\dot E$ is the pulsar spin-down power and $\eta$ is the filling factor. In general, if the pulsar is subsonic in the ambient medium, i.e. $ P_{\rm ext} \geq P_{\rm ram}= 6\times 10^{-10}nv^2_7\;\mathrm {g\;cm^{-1}s^{-2}}$ for a pulsar speed $10^7v_7\;\mathrm{cm\;s^{-1}}$ through a density of $n\,m_p\;\mathrm{cm^{-3}}$, a toroidal shock structure is expected; faster objects produce bow shock nebulae. Indeed, many young pulsars still inside their high pressure supernova remnant birth sites do show such toroidal symmetry. The best-known example is the PWN around the \object{Crab pulsar} as observed by the \emph{CXO} \citep{wei01}. Recently, several relativistic MHD models, e.g.\ \citet{kom03} and \citet{del06}, have shown how such toroidal structure can form if the pulsar wind has a latitudinal variation. \citet{ng04} (hereafter \citetalias{ng04}) developed a fitting procedure to measure the 3D orientation of the pulsar wind torus and applied to a few X-ray observations. While this simple geometrical model does not capture the fine details of the MHD simulations, is does allow one to extract the torus (and hence pulsar spin) orientation from relatively low signal-to-noise ratio (S/N) data. \citetalias{ng04} also gave quantitative estimates for the statistical errors arising from Poisson statistics. However the systematic errors due to unmodeled components such as jets or background were neglected. For bright objects e.g.\ the Crab and Vela pulsars, the S/N is high and such systematic errors dominate. In this study, we attempt to quantify these systematic errors and apply the fitting to more \emph{CXO} PWN observations, thus providing a more comprehensive study. | In conclusion, we have applied the torus fitting technique to more PWN observations in the \emph{Chandra} data archive and characterized the uncertainties in the fits. This study provides a better understanding of the systematic errors, giving quantitative estimates of the measurement uncertainties. We argue that these robust position angle $\Psi$ and inclination $\zeta$ values are particularly useful for comparison with the radio and high energy pulse data. If new observations can fill in more measurements from these energy bands in Table 3, we should be able to make substantial progress in understanding the emission zones and viewing geometries of young pulsars. | 7 | 10 | 0710.4168 |
0710 | 0710.1860_arXiv.txt | We present a comprehensive analysis of structure in the young, embedded cluster, NGC~1333 using members identified with {\it Spitzer} and 2MASS photometry based on their IR-excess emission. In total, 137 members are identified in this way, composed of 39 protostars and 98 more evolved pre-main sequence stars with disks. Of the latter class, four are transition/debris disk candidates. The fraction of exposed pre-main sequence stars with disks is $83\% \pm 11\%$, showing that there is a measurable diskless pre-main sequence population. The sources in each of the Class~I and Class~II evolutionary states are shown to have very different spatial distributions relative to the distribution of the dense gas in their natal cloud. However, the distribution of nearest neighbor spacings among these two groups of sources are found to be quite similar, with a strong peak at spacings of 0.045~pc. Radial and azimuthal density profiles and surface density maps computed from the identified YSOs show that NGC~1333 is elongated and not strongly centrally concentrated, confirming previous claims in the literature. We interpret these new results as signs of a low velocity dispersion, extremely young cluster that is not in virial equilibrium. | Observations of embedded, star-forming clusters and groups show that the stellar distributions are often elongated, clumpy, or both \citep[cf.][]{ll03,gute05,teix06,alle07}, and the structure seems tied to the distribution of dense gas in the clusters' natal molecular clouds. Most clouds have some non-spherical structure, yet current results suggest that the associated clusters have varying degrees of agreement with the cloud's structure depending on how deeply the members are embedded \citep{gute05,teix06}. Given the high frequency of asymmetric structure in clouds, it seems reasonable to assume that the exposed, relatively structureless clusters we observe may have been asymmetric in a previous epoch. With the ejection of the majority of their natal gas and adequate time to migrate from their birth sites, the imprint of the underlying cloud structure could be lost rather quickly in a recently exposed young cluster. Thus the structure we measure in a distribution of young stellar objects (YSO) relative to the dense gas in the associated cloud may be a reasonable proxy for the current dynamical state of an embedded cluster. The launch of {\it Spitzer} has provided a potent new facility for studies of the structure of young and embedded clusters. While some members of these young clusters are diskless pre-main sequence stars (Class~III), the majority of the membership are sources with excess emission at mid-IR wavelengths \citep{hll01}, made up of a mixture of protostars (Class~I), still embedded and accreting from dense spherical envelopes, and the slightly more evolved pre-main sequence stars with circumstellar disks (Class~II). The mid-IR spectral energy distributions (SED) of these objects are dominated by the emission from their dusty circumstellar material, making them easily distinguishable from pure photospheric sources such as unrelated field stars and indistinguishable diskless cluster members. {\it Spitzer}'s sensitivity to mid-IR emission makes it the best tool currently available for identifying {\it and characterizing} YSOs with IR-excess, and that sensitivity is sufficient to detect these sources down to the Hydrogen--burning mass limit for regions within the nearest kiloparsec \citep{gute04}. The complete census of YSOs with disks in a young cluster represents a high-confidence sample of bona fide cluster members, and for the youngest regions, such a sample includes a high fraction of the total number of members \citep{hll01}. Furthermore, by separating the two canonical evolutionary classes of YSOs that {\it Spitzer} so effectively detects, we are able to probe both the recent overall star-forming activity in the region as traced by the stars with disks, and the immediate star formation as represented by the protostars, since this phase is expected to be short-lived relative to the former. Many young clusters are dominated by heavy and spatially variable extinction from dust within their natal molecular cloud environment. Another advantage {\it Spitzer} brings to studies of these regions is that the mid-IR wavelengths targeted by {\it Spitzer}'s imaging instruments (Infrared Array Camera, or IRAC, at 3.6-8.0~$\mu$m and Multiband Imaging Photometer for {\it Spitzer}, or MIPS, at 24-160~$\mu$m) are less affected by extinction from dust in comparison to near-IR or visible wavelengths \citep[e.g.][]{ccm89}. Recent papers have done an excellent job of characterizing the reddening law in the IRAC bandpasses \citep{inde05,flah07,huar07} and have made a first attempt at doing the same for the 24~$\mu$m channel of MIPS \citep{flah07,huar07}. Given the wealth of information {\it Spitzer} can bring to the study of embedded clusters, we have surveyed over thirty clustered star-forming regions within the nearest kiloparsec through the Guaranteed Time Observations (GTO) program for the IRAC and MIPS instrument teams. With these data, we can not only provide a full census of the YSOs with IR-excess in each region, but also perform a detailed examination and comparative analysis of structure in young, embedded clusters. NGC~1333 has been a popular target for observations of deeply embedded protostars via radio wavelengths due to the presence of an unprecendented number of molecular outflows \citep[e.g.][]{ks00} associated with several bright IRAS sources considered Class~0 protostars \citep{jenn87}, all in relatively close proximity to the Sun \citep[250~pc,][]{enoc06}. Many of the outflows are traced by shock-induced emission, a clear sign that they are indeed affecting the local, quiescent cloud material \citep[e.g.][]{wala05}. Because of this, some studies have claimed that the NGC~1333 dense molecular cloud core is in the process of being destroyed by influence from outflows \citep[e.g.][]{wari96,sk01,quil05}. In addition to the numerous protostars, there is a cluster of pre-main sequence stars identified by number counts analysis in the near-IR \citep{svs76,asr94,lal96,wilk04}. \citet{lal96} suggested that the distribution was well-described as a ``double cluster'', having two distinct surface density maxima. Here we present a {\it Spitzer} IRAC and MIPS imaging and photometric analysis of the NGC~1333 young cluster, one of the most nearby large membership ($N>100$) clusters in the {\it Spitzer} Young Cluster Survey. We achieve a near-complete census of the cluster membership that possesses circumstellar material, significantly surpassing the sensitivity of ground-based mid-IR surveys \citep[cf.][and references therein.]{rebu03}. In addition, we statistically infer the population of pre-main sequence stars that lack disks, enabling an estimate of the fraction of members with disks. Finally, using the identified YSOs, we apply both established and recently developed methods for characterizing the structure of this cluster, in specific reference to previous claims of structure found in the literature. | We have employed the mid-IR sensitivity of {\it Spitzer} to achieve a census of the YSO members of the NGC~1333 embedded cluster that was previously unattainable. Furthermore, we have shown that the sources identified in our {\it Spitzer} census represent a large fraction of the total cluster membership (83\%). With this penetrating view of the cluster, we have performed several measurements of different aspects of the structure of the cluster that are not confused by field star contamination and are less biased by extinction effects than previous studies. We confirm the double-peaked surface density morphology of the cluster reported in previous work \citep{lal96}, with the caveat that this morphology is traced only by the more evolved Class~II population. The two main density peaks are located in local minima of the gas density distribution. In contrast, the protostars trace the dense gas in this region closely, as do a fraction of the Class~II. That gas distribution appears to be a network of filamentary structures, connecting the two density peaks into a single, albeit complex, structure. Despite the difference in spatial distributions, however, the Class~II and protostars have similarities in their nearest neighbor distance distributions, particularly in the location of the peak at spacings of 0.045~pc. These two evolutionary states are expected to differ in duration by an order of magnitude, thus a low overall velocity dispersion in the cluster and a very young Class~II population are needed to account for the similarity remaining in the spacings among the two populations. A further implication of the lack of evidence for dynamical evolution of the cluster is the need for the dispersal of the molecular gas by the Class~II sources on very short timescales. Given that the more evolved YSOs have not moved significantly from their birthsites, regions dominated by more evolved sources where gas densities are preferentially low, such as the double-peaks in the stellar distribution here, were once locations of high gas density. Thus it seems plausible that groups of low-mass stars are able to disrupt their dense natal gas {\it locally}. In this sense, we concur with previous work that has suggested that the active outflows in NGC~1333 are destroying the cloud \citep{sk01,quil05}. With a virtual lack of low or high mass stars and dense gas in the center of the ring-like structure of the cloud and cluster core, we argue that this is in fact a primordial structure and not outflow-evacuated. Considering all sources regardless of evolutionary class, the cluster is clearly elongated, an expected result given the double-peaked nature of the stellar surface density distribution and embedded nature of the cluster as a whole \citep{gute05}. We have presented radial density profiles measured via several different methods, and all suggest a roughly uniform density distribution within a 0.3~pc radius, with a steep decline ($\alpha=-3.3$) beyond. The flat central distribution and sharp radial decline in the surface density profile at radii larger than 0.3~pc also suggests a rather limited amount of dynamical interaction in this cluster. If we analyze this within the framework of dynamics-generated structures like those of King or EFF models, we have to choose extreme fitting parameters just to get marginal agreement with the measured profiles. The radial profile we have measured here has two clear regimes were a simple power law matchs the density profile well, and the sharp knee transition between them is poorly matched by either the King or EFF profiles. Furthermore, given the asymmetry of the NGC~1333 cluster and the presense of a considerable amount of structured, dense gas, the dynamical information inferred from either of these models is unlikely to be accurate. Cloud geometry is the most tenable cause for the structure we have observed. As such, the relative motions of the stars must be fairly slow and the Class~II population must be quite young in order to have preserved that structure into the Class~II evolutionary phase. | 7 | 10 | 0710.1860 |
0710 | 0710.3054_arXiv.txt | We report on the spatial relationship between solar flares and coronal mass ejections (CMEs) observed during 1996-2005 inclusive. We identified 496 flare-CME pairs considering limb flares (distance from central meridian $\ge 45^\circ$) with soft X-ray flare size $\ge$ C3 level. The CMEs were detected by the Large Angle and Spectrometric Coronagraph (LASCO) on board the {\it Solar and Heliospheric Observatory} ({\it SOHO}). We investigated the flare positions with respect to the CME span for the events with X-class, M-class, and C-class flares separately. It is found that the most frequent flare site is at the center of the CME span for all the three classes, but that frequency is different for the different classes. Many X-class flares often lie at the center of the associated CME, while C-class flares widely spread to the outside of the CME span. The former is different from previous studies, which concluded that no preferred flare site exists. We compared our result with the previous studies and conclude that the long-term LASCO observation enabled us to obtain the detailed spatial relation between flares and CMEs. Our finding calls for a closer flare-CME relationship and supports eruption models typified by the CSHKP magnetic reconnection model. | A solar flare is sudden flash of electromagnetic radiation (suggesting plasma heating) in the solar atmosphere, and a coronal mass ejection (CME) is an eruption of the atmospheric plasma into interplanetary space. Both phenomena are thought to be different manifestations of the same process which releases magnetic free energy stored in the solar atmosphere. The spatial relation between flares and CMEs contains information on the magnetic field configurations involved in the eruptive process and hence is important for modeling them. Many flare-CME models are based on the CSHKP (Carmichael, Sturrock, Hirayama, Kopp \& Pneuman) magnetic reconnection model. The model requires that a flare occurs just underneath of an erupting filament which eventually becomes the core of the CME associated with the flare. Normally the core corresponds to the center of the CME, thus the CSHKP model requires that the flare occurs near the center of the CME span. Full-scale studies on the flare-CME relationship started in the 70s and 80s with the CME observations obtained by the {\it Solwind} coronagraph on board {\it P78-1} and the Coronagraph/Polarimeter telescope on board the {\it Solar Maximum Mission} ({\it SMM}). \citet{harri86} carried out a detailed analysis of three flare-CME events observed by {\it SMM} and reported that flares occurred near one foot of an X-ray arch, which is supposed to become a CME. He also analyzed 48 flare-CME events observed by {\it SMM} and {\it Solwind} and reported that many flares occurred near one leg of the associated CMEs. This result, called the flare-ejection asymmetry, is inconsistent with the CSHKP flare-CME model. \citet{kahle89} examined 35 events observed by the {\it Solwind} and reported that flare positions did not peak neither at the center nor at one leg of the CMEs. They concurred with Harrison at the point that the observations do not match with the CSHKP model, while disagreeing with the result that flares are likely to occur at one leg of CMEs. They pointed out that the parameter employed by Harrison was biased, and concluded that both observations are compatible with the fact that there is no preferred flare site with respect to the CME span. It should be noted that the two studies applied different criteria for the event selection. Harrison did not apply any criteria on flare X-ray intensity, flare location, and CME span, while Kahler et al. used only strong limb flares ($\ge$ M1 level; central meridian distance (CMD) $\ge 40^\circ$) and wide CMEs (angular span $\ge 40^\circ$). Different criteria might produce different spatial distributions, but the results in both the studies were inconsistent with the schematic view of the CSHKP type flare-CME model. The Large Angle and Spectrometric Coronagraph \citep[LASCO;][]{bruec95} on board the {\it Solar and Heliospheric Observatory} ({\it SOHO}) mission has observed more than 11,000 CMEs from 1996, which provides a great opportunity to investigate the flare-CME relationship. \citet{harri06} reviewed several flare-CME studies and stated that "the pre-SOHO conclusions about relative flare-CME locations and asymmetry are consistent with many recent studies." However, systematic statistical study is needed before reaching a firm conclusion. In this paper we revisit this issue using the large CME data obtained by {\it SOHO} LASCO. \begin{figure*} \epsscale{0.90} \plotone{f1.eps} \caption{Three CMEs observed by {\it SOHO} LASCO to illustrate the measurement of CME span. The top row shows direct images used to measure the main CME body, and the bottom row shows corresponding running difference images used to measure the whole CMEs. $\phi_1$ and $\phi_2$ indicate the PAs of side edges of the main CME body, and $\phi_A$ and $\phi_B$ indicate those of the whole CME. Arrows point to the position of the flares associated with the CMEs.} \end{figure*} | 7 | 10 | 0710.3054 |
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0710 | 0710.1098_arXiv.txt | % A planetary transit produces both a photometric signal and a spectroscopic signal. Precise observations of the transit light curve reveal the planetary radius and allow a search for timing anomalies caused by satellites or additional planets. Precise measurements of the stellar Doppler shift throughout a transit (the Rossiter-McLaughlin effect) place a lower bound on the stellar obliquity, which may be indicative of the planet's migration history. I review recent results of the Transit Light Curve project, and of a parallel effort to measure the Rossiter effect for many of the known transiting planets. | % I have great admiration for the people who discover transiting planets. Identifying the candidate transit signals from among a hundred thousand light curves, and flushing out the numerous astrophysical false positives, are impressive feats. This article, however, is not about transit discovery, but rather about the next step: performing high-precision photometry and spectroscopy of exoplanetary transits. The goal of this step is to determine the planetary and stellar properties well enough to allow for meaningful comparisons with the familiar properties of the Solar system, and to inform our theories of planet formation. The most immediate result of transit photometry is a measurement of the planetary radius. In combination with the planetary mass, which can be inferred from the Doppler orbit of the star, these data give the first clues about the composition, interior structure, and atmospheric energy balance of the planet. An accurate radius is also needed to interpret the results of other observations, such as the detection of thermal emission or reflected light based on secondary-transit photometry. The timings of the transits can be used to refine the measurement of the orbital period and search for additional bodies in the system. In \S~2, I describe the Transit Light Curve (TLC) project, an effort to gather high-precision photometry during exoplanetary transits. The most prominent spectroscopic signal during a transit is the Rossiter-McLaughlin (RM) effect. This effect is an anomalous Doppler shift that arises from stellar rotation. Measuring this effect allows one to assess the alignment between the planetary orbital axis and the stellar spin axis, a fundamental system property that provides clues about the process of planet migration. I describe some recent measurements of the RM effect in \S~3. \begin{figure}[!h] \plotone{Lightcurves.eps} \caption{{\bf The Transit Light Curve project.} Each panel shows time-binned photometry from a campaign (2-5 separate transits) for a particular planet. Most of the data are $z$ band observations with the FLWO 1.2m telescope and Keplercam detector.} \end{figure} | 7 | 10 | 0710.1098 |
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0710 | 0710.0866_arXiv.txt | We discuss oxygen and iron abundance patterns in K and M red-giant members of the Galactic bulge and in the young and massive M-type stars inhabiting the very center of the Milky Way. The abundance results from the different bulge studies in the literature, both in the optical and the infrared, indicate that the [O/Fe]-[Fe/H] relation in the bulge does not follow the disk relation, with [O/Fe] values falling above those of the disk. Based on these elevated values of [O/Fe] extending to large Fe abundances, it is suggested that the bulge underwent a rapid chemical enrichment with perhaps a top-heavy initial mass function. The Galactic Center stars reveal a nearly uniform and slightly elevated (relative to solar) iron abundance for a studied sample which is composed of 10 red giants and supergiants. Perhaps of more significance is the fact that the young Galactic Center M-type stars show abundance patterns that are reminiscent of those observed for the bulge population and contain enhanced abundance ratios of $\alpha$-elements relative to either the Sun or Milky Way disk at near-solar metallicities. | Abundance patterns in different populations in the Milky Way can shed light on Galaxy formation and its chemical evolution. Studies carried out over the last decade have provided accurate abundance patterns for stars in the Milky Way disk and halo so that these populations are now fairly-well mapped. There is significantly less information, however, on the Galactic bulge population and, in particular the Galactic Center, due to difficulties associated with heavy extinction. The stellar content belonging to the bulge population is old, with an age of 10--12 Gyr (e.g. Zoccali et al. 2003) and chemical abundance studies in both low and high-resolution find that the bulge is overall metal-rich with a large abundance spread that spans $\sim$1.5 dex (e.g. Fulbright et al. 2006). The Galactic Center, on the other hand, contains a significant young population, including many luminous and massive stars. Those known to date are concentrated in three clusters within 60 pc of the Galactic Center: the Central Cluster, the Arches Cluster and the Quintuplet Cluster. The youngest stars in these clusters were formed recently with ages ranging roughly between 1 -- 9 Myr. In this contribution we summarize the abundance results obtained recently for the bulge and Galactic center and briefly discuss their implications in light of Galactic chemical evolution. | Abundance patterns are discussed in two distinct populations: $\sim$ 12 Gyr old red-giants from the bulge and the young and massive M-type giants and supergiants residing mostly within 2 pcs of the central Black Hole. The abundance results indicate that the iron abundance distribution for the admittedly small sample of Galactic Center targets studied so far shows very little abundance spread in contrast to the large metallicity spread found for the old bulge population. Most importantly, both populations show enhancements in the $\alpha$-element abundances relative to the Galactic disk trend that might be explained invoking a top-heavy IMF. | 7 | 10 | 0710.0866 |
0710 | 0710.2114_arXiv.txt | { We present new determination of the birth rate of AXPs and SGRS and their associated SNRs. We find a high birth rate of $1/(500\ {\rm yr})$ to $1/(1000\ {\rm yr})$ for AXPs/SGRs and their associated SNRs. These high rates suggest that all massive stars (greater than $\sim (23$-$32) M_{\odot}$) give rise to remnants with magnetar-like fields. Observations indicate a limited fraction of high magnetic fields in these progenitors thus our study necessarily implies magnetic field amplification. Dynamo mechanisms during the birth of the neutron stars require spin rates much faster than either observations or theory indicate. Here, we propose that neutron stars form with normal ($\sim 10^{12}$ G) magnetic fields, which are then amplified to $10^{14}$-$10^{15}$ G after a delay of hundreds of years. The amplification is speculated to be a consequence of color ferromagnetism and to occur after the neutron star core reaches quark-deconfinement density. This delayed amplification alleviates many difficulties in interpreting simultaneously the high birth rate, high magnetic fields, and state of isolation of AXPs/SGRs and their link to massive stars. | Early studies of association of Anomalous X-ray Pulsars (AXPs) with supernova remnants (SNRs) suggested that 5\% of core-collapse SN results in AXPs (Gaensler et al. 1999). This was based on 3 SNR associations out of a total of 6 AXPs. Since then evidence has mounted that AXPs and soft gamma-ray repeaters (SGRs) are the same type of objects (Gavriil et al. 2002) and more AXPs, SGRs and associated SNRs have been identified. Thus it it timely to revisit the issue of AXPs/SGRs birthrates. In this study we present an updated investigation of the birth rate of AXPs/SGRs and in addition, for the first time, the birth rate of associated SNRs is given. Since AXPs/SGRs ages rely on spin-down age estimates whereas SNRs ages are based on shock expansion models, this constitutes two independent estimates for birth rates. Both samples yield a high birth rate for AXPs/SGRs\footnote{An independent study by Gill\&Heyl (2007), based on a population synthesis of AXPs detected in the ROSAT All-Sky Survey, yields a birth rate of $\sim 0.22$ per century.} of $(1/5)$-$(1/10)$ of all core-collapse SNe, higher than previously appreciated. This high frequency of occurrence of AXPs/SGRs brings into focus issues related to the origin of the strong magnetic fields which we address here. This paper is presented as follows: \S 2 describes the methods and presents the birth rate results, and \S 3 discusses the implications. Our model, based on a delayed amplification of magnetic field, is presented in \S 4 before concluding in \S 5. | Our study of the birth rate of AXPs and SGRS and their associated SNRs suggest that about $1/5$ to $1/10$ of all core-collapse SN lead to AXPs/SGRs. These high rates suggest that all massive stars (greater than $ M_{\rm low}$) give rise to remnants with magnetar-like fields. This raises these issues: (i) how do all progenitors with $M\ge M_{\rm low}$ generate $>10^{14}$ G fields in their compact remnants?; (ii) why is there a dichotomy in magnetic field strength between compact remnants from progenitors with mass greater than $M_{\rm low}$ (i.e. $B\sim 10^{14}$ G) and those with mass less than $\sim M_{\rm low}$ ($B\sim 10^{12}$ G); (iii) why are all AXPs/SGRs isolated while many progenitors with $M>M_{\rm low}$ are in binaries? In this study, we introduce the notion of delayed magnetic field amplification to resolve these issues. We propose that neutron stars from progenitor masses $M> 9M_{\odot}$ are born with normal ($\sim 10^{12}$ G) magnetic fields. A neutron star from a progenitor with an approximate mass range $M_{\rm low}< M <60 M_{\odot}$ will experience an explosive transition to a quark star (the QN) in which its magnetic field is amplified to $10^{14}$-$10^{15}$ G by color ferromagnetism (Iwazaki 2005). The second explosion (QN) and related mass loss helps to reduce the surviving compact binary fraction thus explaining the state of isolation of AXPs/SGRs. The transition occurs with a delay of several hundred years (Staff et al. 2006). This delayed amplification alleviates many difficulties in interpreting simultaneously the high birth rate and high magnetic fields of AXPs/SGRs and their link to massive stars. | 7 | 10 | 0710.2114 |
0710 | 0710.0920_arXiv.txt | Electric currents $j$ flow along the open magnetic field lines from the polar caps of neutron stars. Activity of a polar cap depends on the ratio $\alph=j/c\rhoGJ$, where $\rhoGJ$ is the corotation charge density. The customary assumption $\alph\approx 1$ is not supported by recent simulations of pulsar magnetospheres and we study polar caps with arbitrary $\alph$. We argue that no significant activity is generated on field lines with $0<\alph<1$. Charges are extracted from the star and flow along such field lines with low energies. By contrast, if $\alph>1$ or $\alph<0$, a high voltage is generated, leading to unsteady $e^\pm$ discharge on a scale-height smaller than the size of the polar cap. The discharge can power observed pulsars. Voltage fluctuations in the discharge imply unsteady twisting of the open flux tube and generation of Alfv\'en waves. These waves are ducted along the tube and converted to electromagnetic waves, providing a new mechanism for pulsar radiation. | Corotation of a plasma magnetosphere is impossible beyond the light cylinder of a star, and magnetic field lines that extend beyond this cylinder are twisted, $\nabla\times\bB\neq 0$. Thus, currents $\bj_B=(c/4\pi)\nabla\times\bB$ are induced in the open magnetic flux tubes that connect the star (its ``polar caps'') to the light cylinder (Sturrock 1971). These currents are approximately force-free and flow along the magnetic field $\bB$. A basic question of pulsar theory is what voltage develops along the open tube to maintain these currents. It determines the dissipated power and $e^\pm$ creation that feeds the observed activity of pulsars. The customary pulsar model assumes that the electric current $\jB$ extracted from the polar cap nearly matches $c\rhoGJ$, where $\rhoGJ=-{\bf \Omega}\cdot\bB/2\pi c$ is the corotation charge density (Goldreich \& Julian 1969). The deviation of current from $c\rhoGJ$ was calculated as an eigen value of an electrostatic problem and found to be small (Arons \& Scharlemann 1979). This is in conflict with recent global models of pulsar magnetospheres, which report $|\jB-c\rhoGJ|\sim\jB$ (e.g. Contopoulos et al. 1999; Spitkovsky 2006; Timokhin 2006; see Arons 2008 for a review). A significant mismatch between $\jB$ and $c\rhoGJ$ can be expected on general grounds (Kennel et al. 1979). Currents $\jB$ are determined by the magnetic-field twisting near the light cylinder, while $\rhoGJ$ is a local quantity at the polar cap that is practically independent of $\jB$. The open tube is surrounded by the grounded closed magnetosphere\footnote{ The closed magnetosphere with $\jB=0$ is expected to have $\rho=\rhoGJ$ and $\Epar\approx 0$ (e.g. Arons 1979). } and may be thought of as a waveguide, filled with magnetized (1D) plasma. Compared to usual plasma-filled waveguides, it has two special features: (1) Current $\jB$ is imposed on the tube. The twisted tube extends into the star, which is a good conductor and maintains $j_B$. ${\rm sign}(\jB)=\pm 1$ corresponds to $\pm$ charge flowing outward along the magnetic field lines. (2) Vacuum has effective charge density $\rho_0=-\rhoGJ$ as Gauss law in the rotating frame reads $\nabla\cdot\bE=4\pi(\rho-\rhoGJ)$. The key dimensionless parameter (which can vary along and across the tube) is \be \alph=\frac{\jB}{c\rhoGJ}. \ee In this Letter, we discuss basic properties of the polar-cap accelerator with arbitrary $\alph$. First, we discuss what happens without $e^\pm$ creation: \S~\ref{sec:slab} studies how the current is extracted from the polar cap of a radius $\rpc$ and flows at small heights $z\ll\rpc$ (this region is called ``slab zone'' below), and \S~\ref{sec:tube} discusses how the flow extends to the region $z>\rpc$ (``thin-tube zone''). We argue that the accelerator is inefficient if $0<\alph<1$. The value of $\alph$ depends largely on the angle $\chi$ between ${\bf\Omega}$ and $\bB$ at the polar cap. For aligned dipole rotators ($\chi=0$), magnetospheric models predict $\alph<1$ everywhere on the polar cap and $\alpha<0$ near its edge (see Fig.~4 and 5 in Timokhin 2006). For orthogonal rotator ($\chi\approx\pi/2$), $|\alph|\gg 1$ throughout most of the polar cap. Generally, the polar cap has three regions where $\alph>1$, $0<\alph<1$, and $\alph<0$. We propose that pulsar activity originates in the polar-cap regions where $\alph^{-1}<1$ (i.e. $\alph>1$ or $\alph<0$) as a high voltage is generated in these regions. We emphasize that the voltage is generated because $\nabla\times\bB\neq 0$, not because $\rhoGJ\neq 0$. The accelerator works as well if $\rhoGJ=0$ (i.e. if $\bf\Omega\perp\bB$ at the polar cap), which corresponds to $\alph\rightarrow \pm\infty$. The accelerator height $h\simlt\rpc$ is regulated by unsteady $e^\pm$ discharges (\S~\ref{sec:discharge}). \S~\ref{sec:Alfven} describes a mechanism for radio emission from unsteady discharges. | 7 | 10 | 0710.0920 |
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0710 | 0710.5734_arXiv.txt | Big Bang nucleosynthesis requires a fine balance between equations of state for photons and relativistic fermions. Several corrections to equation of state parameters arise from classical and quantum physics, which are derived here from a canonical perspective. In particular, loop quantum gravity allows one to compute quantum gravity corrections for Maxwell and Dirac fields. Although the classical actions are very different, quantum corrections to the equation of state are remarkably similar. To lowest order, these corrections take the form of an overall expansion-dependent multiplicative factor in the total density. We use these results, along with the predictions of Big Bang nucleosynthesis, to place bounds on these corrections and especially the patch size of discrete quantum gravity states. | \label{sec:INTRODUCTION} Much of cosmology is well-described by a space-time near a spatially isotropic Friedmann--Robertson--Walker models with line elements \begin{equation} \md s^2= -\md\tau^2+a(\tau)^2\left(\frac{\md r^2}{1-kr^2}+r^2(\md\vartheta^2+ \sin^2\vartheta\md\varphi^2)\right) \end{equation} where $k=0$ or $\pm 1$, sourced by perfect fluids with equations of state $P=w\rho$. Such an equation of state relates the matter pressure $P$ to its energy density $\rho$ and captures the thermodynamical properties in a form relevant for isotropic space-times in general relativity. Often, one can assume the equation of state parameter $w$ to be constant during successive phases of the universe evolution, with sharp jumps between different phases such as $w=-1$ during inflation, followed by $w=\frac{1}{3}$ during radiation domination and $w=1$ during matter domination. Observationally relevant details can depend on the precise values of $w$ at a given stage, in particular if one uses an effective value describing a mixture of different matter components. For instance, during big bang nucleosynthesis one is in a radiation dominated phase mainly described by photons and relativistic fermions. Photons, according to Maxwell theory, have an exact equation of state parameter $w=\frac{1}{3}$ as a consequence of conformal invariance of the equations of motion (such that the stress-energy tensor is trace-free). For fermions the general equation of state is more complicated and non-linear, but can in relativistic regimes be approximately given by the same value $w=\frac{1}{3}$ as for photons. In contrast to the case of Maxwell theory, however, there is no strict symmetry such as conformal invariance which would prevent $w$ to take a different value. It is one of the main objectives of the present paper to discuss possible corrections to this value. For big bang nucleosynthesis, it turns out, the balance between fermions and photons is quite sensitive. In fact, different values for the equation of state parameters might even be preferred phenomenologically \cite{FermionBoson}. One possible reason for different equations of state could be different coupling constants of bosons and fermions to gravity, for which currently no underlying mechanism is known. In this paper we will explore the possibility whether quantum gravitational corrections to the equations of state can produce sufficiently different values for the equation of state parameters. In fact, since the fields are governed by different actions, one generally expects different, though small, correction terms which can be of significance in a delicate balance. Note that we are not discussing ordinary quantum corrections of quantum fields on a classical background. Those are expected to be similar for fermions and radiation in relativistic regimes. We rather deal with quantum gravity corrections in the coupling of the fields to the space-time metric, about which much less is known a priori. Thus, different proposals of quantum gravity may differ at this stage, providing possible tests. An approach where quantum gravitational corrections can be computed is loop quantum cosmology \cite{LivRev}, which specializes loop quantum gravity \cite{Rov,ALRev,ThomasRev} to cosmological regimes. In such a canonical quantization of gravity, equations of state must be computed from matter Hamiltonians rather than covariant stress-energy tensors. Quantum corrections to the underlying Hamiltonian then imply corrections in the equation of state. This program was carried out for the Maxwell Hamiltonian in \cite{MaxwellEOS}, and is done here for Dirac fermions. There are several differences between the treatment of fermions and other fields, which from the gravitational point of view are mainly related to the fact that fermions, in a first order formulation, also couple to torsion and not just the curvature of space-time. After describing the classical derivation of equations of state as well as steps of a loop quantization and its correction terms, we use big bang nucleosynthesis constraints to see how sensitively we can bound quantum gravity parameters. We have aimed to make the paper nearly self-contained and included some of the technical details. Secs.~\ref{sec:canonical Formulation} on the canonical formulation of fermions, \ref{sec:MODIFICATIONS} on quantum corrections from loop quantum gravity and \ref{sec:BBN} on the analysis of big bang nucleosynthesis can, however, be read largely independently of each other by readers only interested in some of the aspects covered here. We will start with general remarks on the physics underlying the problem. | Big bang nucleosynthesis is a highly relativistic regime which, to a good approximation, implies identical equations of state for fermions and photons. There are, however, corrections to the simple equation of state $w=\frac{1}{3}$ for fermions even classically. One observation made here is that the interaction term derived in \cite{FermionImmirzi} leads to such a correction and might be more constrained by nucleosynthesis than through standard particle experiments \cite{FermionTorsion}. We have not analyzed this further here because more details of the behavior of the fermion current would be required. A second source of corrections arises from quantum gravity. Remarkably, while quantum gravity effects on an isotropic background do correct the equations of state, they do so equally for photons and relativistic fermions. Initially, this is not expected for both types of fields due to their very different actions. Thus, quantum gravity effects do not spoil the detailed balance required for the scenario to work and bounds from big bang nucleosynthesis obtained so far are not strong. But there are interesting limits for the primary parameter, the patch size of a quantum gravity state. It is dimensionally expected to be proportional to the Planck length $\ell_{\rm P}$ but could be larger. In fact, current bounds derived here already rule out a patch size of exactly the elementary allowed value in loop quantum gravity. With more precise estimates, these bounds can be improved further. We have made use of quantum gravity corrections in a form which does not distinguish fermions from radiation. Although the most natural implementation, quite unexpectedly, provides equal corrections as shown here, there are several possibilities for differences which suggest several further investigations. Small deviations in the equations of state and thus energy densities of fermions and radiation are possible. First, there are always quantization ambiguities, and so far we tacitly assumed that the same basic quantization choice is made for the Maxwell and Dirac Hamiltonians. Such ambiguity parameters can be explicitly included in specific formulas for correction functions; see e.g.\ \cite{Ambig,ICGC,QuantCorrPert}. Independent consistency conditions for the quantization may at some point require one to use different quantizations for both types of fields, resulting in different quantum corrections and different energy densities. Such conditions can be derived from an analysis of anomaly-freedom of the Maxwell field and fermions coupled to gravity, which is currently in progress. As shown here, if this is the case it will become testable in scenarios sensitive to the behavior of energy density such as big bang nucleosynthesis. Moreover, assuming the same quantization parameters leads to identical quantum corrections for photons and fermions only on isotropic backgrounds. Small-scale anisotropies have different effects on both types of fields and can thus also be probed through their implications on the equation of state. For this, it will be important to estimate more precisely the typical size of corrections, which is not easy since it requires details of the quantum state of geometry. The crucial ingredient is again the patch size of underlying lattice states. On the other hand, taking a phenomenological point of view allows one to estimate ranges for patch sizes which would leave one in agreement with big bang nucleosynthesis constraints. Interestingly, corrections studied here provide upper bounds to the patch size, and other corrections from quantum gravity are expected to result in lower bounds. A finite window thus results, which can be shrunk with future improvements in observations. | 7 | 10 | 0710.5734 |
0710 | 0710.3732_arXiv.txt | % Edge-on spiral galaxies offer a unique perspective on disks. One can accurately determine the height distribution of stars and ISM and the line-of-sight integration allows for the study of faint structures. The Spitzer IRAC camera is an ideal instrument to study both the ISM and stellar structure in nearby galaxies; two of its channels trace the old stellar disk with little extinction and the 8 micron channel is dominated by the smallest dust grains (Polycyclic Aromatic Hydrocarbons, PAHs). \cite{Dalcanton04} probed the link between the appearance of dust lanes and the disk stability. In a sample of bulge-less disks they show how in massive disks the ISM collapses into the characteristic thin dust lane. Less massive disks are gravitationally stable and their dust morphology is fractured. The transition occurs at 120 km/s for bulgeless disks. Here we report on our results of our Spitzer/IRAC survey of nearby edge-on spirals and its first results on the NIR Tully-Fischer relation, and ISM disk stability. | For 32 edge-on galaxies, spanning Hubble type and mass, we fit the edge-on infinite disk model by \cite{vdKruit81a} on the IRAC mocaics; the stellar dominated 4.5 $\mu$m and the PAH emission at 8 $\mu$m, with the stellar contribution subtracted \citep[][]{Pahre04a}. The disk's total luminosity is inferred from the fitted model: $L_{disk} = 2 \pi h^2 \mu_0$, with $h$ the scale-length and $L_0$ the face-on central surface brightness \citep[][]{Kregel02}. Figure \ref{f:tf} shows the inferred Tully-Fischer relation for these disks for the stars (4.5 $\mu$m). Notably, the slope ($\alpha$) is 3.5, similar to what \cite{Meyer06a} found but contrary for the increasing trend of slope with redder filters. The effects of age and metallicity of the stellar population become independent and opposite effects on the color-M/L relation in NIR \citep[See][Fig. 2d]{BelldeJong}. Hence, the shallower slope in the IRAC stellar channels, could be the metallicity effect starting to dominate. | 7 | 10 | 0710.3732 |
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0710 | 0710.4262_arXiv.txt | The accuracy of position measurements on stellar targets with the future Space Interferometry Mission (SIM) will be limited not only by photon noise and by the properties of the instrument (design, stability, etc.) and the overall measurement program (observing strategy, reduction methods, etc.), but also by the presence of other ``confusing'' stars in the field of view (FOV). We use a simple ``phasor'' model as an aid to understanding the main effects of this ``confusion bias'' in single observations with SIM. This analytic model has been implemented numerically in a computer code and applied to a selection of typical SIM target fields drawn from some of the Key Projects already accepted for the Mission. We expect that less than 1\% of all SIM targets will be vulnerable to confusion bias; we show that for the present SIM design, confusion may be a concern if the surface density of field stars exceeds 0.4 star/arcsec$^2$. We have developed a software tool as an aid to ascertaining the possible presence of confusion bias in single observations of any arbitrary field. Some \textit{a priori} knowledge of the locations and spectral energy distributions of the few brightest stars in the FOV is helpful in establishing the possible presence of confusion bias, but the information is in general not likely to be available with sufficient accuracy to permit its removal. We discuss several ways of reducing the likelihood of confusion bias in crowded fields. Finally, several limitations of the present semi-analytic approach are reviewed, and their effects on the present results are estimated. The simple model presented here provides a good physical understanding of how confusion arises in a single SIM observation, and has sufficient precision to establish the likelihood of a bias in most cases. We close this paper with a list of suggestions for future work on this subject. | \label{intro} The Space Interferometry Mission (SIM\footnote{also currently called SIM--PlanetQuest}) is being designed by NASA's Jet Propulsion Laboratory to provide a facility-class instrument for measuring the positions and proper motions of stars at optical wavelengths with micro-arc-second (\muas) precision. This represents an improvement by several orders of magnitude over the precision of all existing astrometric instruments. For faint sources (V $\gtrsim 15$ mag.), SIM will also be more than a factor 10 better than any other future planned space mission, and will therefore uniquely permit many new classes of problems to be addressed. Such problems include: the direct measurement (for the first time) of the masses of earth-like planets in orbit around nearby stars; determining the distances to stars by direct triangulation over the whole Galaxy and out to the Magellanic Clouds; measuring the transverse motions of galaxies in the Local Group; and, establishing the shape of the dark matter distribution in the Galaxy. A description of the instrument and its current science program is available at the JPL/SIM web site.\footnote{See {\it http://planetquest.jpl.nasa.gov/SIM/sim\_index.cfm} for descriptions of the current set of Key Projects. Almost half of the available 5-year Mission time is still unallocated.} The mission is now at the end of the detailed design phase. After more than 15 years of development, all the major technical questions have been answered. New devices have been invented in order to provide metrology internal to the spacecraft at a level of a few tens of picometers, a fraction of the inter-atomic distance in a molecule of oxygen. The next major step is to begin construction of the instrument. SIM will be the second optical interferometer in space devoted to astrometry, following the Fine Guidance Sensors (FGS) on the Hubble Space Telescope (HST). However, SIM is a Michelson interferometer using separated collectors, quite different from the filled-aperture white-light shearing interferometer design of the FGS. SIM includes three long-baseline interferometers housed on a common truss, each formed by two $\approx 0.3$ m apertures which compress their light beams and guide them through delay lines to beam combiners. During operation, two interferometers are used for precision guiding of the spacecraft, and the third views the target of interest. Data are then accumulated by tracking the target until enough photons have been recorded to achieve the particular science goal. The precision with which astrometry can be done on a specific stellar target with SIM will be limited by photon noise, the design and stability of the instrument, and by the data calibration and processing. We have some control over the instrument properties, which depend on the specific choices made when implementing the design in hardware. Also, the operation and calibration of the instrument can be optimized so as to maximize the achievable precision. However, there is another source of error which may be present, and which is largely out of our control; it does not reduce the ultimate \textit{precision} with which a given \textit{single measurement} can be made with SIM, but it may reduce the final \textit{accuracy} of that measurement. This source of error arises because of the presence of other stars in the SIM field of view (FOV) which can ``confuse'' any single observation made on the target star. The light from these extraneous stars perturbs the measurement, so that the measured target position can differ from the true position. The difference is a \textit{bias} which can reach a level of many times the single-measurement precision estimated from the instrument parameters alone. It is this \textit{single-measurement confusion bias} which concerns us in this paper. It is important to emphasize that the final \textit{accuracy} of the astrometric parameters (position, parallax, and proper motion) determined by SIM for any given target star will be a result of carrying out a complex program of several single measurements on that target, plus repeated measurements on many other stars for the determination of calibration and baseline orientation parameters \citep{bod04,miltur03}. Since we are concerned only with possible bias in a \textit{single measurement}, the details of the entire observing program are not directly relevant here; however, it also means that we can not quantify the consequences of confusion bias to the final accuracy with which the ``end-of-mission'' astrometric parameters can be obtained on any specific target. It is clear that the effects of single-measurement confusion bias will generally diminish as more observations are combined. But this also means that projects which involve only a few observations of a target (e.g.\ ``targets of opportunity'', single parallax measurements on nearby bright stars, etc.) may have a greater susceptibility to confusion bias. Specific aspects of confusion in astrometric measurements with SIM have been considered by several authors in the recent past. \citet{dalalgriest01} showed that the characteristic response of SIM's fixed-baseline interferometer as a function of wavenumber and delay can be used to refine a model of the distribution of confusing stars in the FOV. This model can then be used to correct the measured position of the target of interest, and in many cases where the level of confusion is not too great the astrometric accuracy can approach the measurement precision. Dalal and Griest successfully applied their method to models of confused fields in the LMC in which $\approx 16$ faint stars are scattered over the FOV around the $\approx 19$ magnitude target star. Photon noise is included in these models. These authors then go on to consider the additional complication if one of the stars in the FOV changes brightness owing to a micro-lensing event, and show that an extension of their fitting algorithm to include the precision photometry provided by SIM's detectors permits even this apparently-intractable case to be handled almost as well. However, their method fails when the angular separation between any pair of sources in the FOV (as projected on the interferometer baseline) corresponds to a delay difference of $\lesssim 2$ coherence lengths for the full bandpass of the detection system. This is a projected separation of 25 milli-arcseconds (mas). Indeed, this is a \textit{general limit} for SIM observations. Our approach is somewhat simpler than that of Dalal and Griest, and yields some improvement in the minimum angular separation which can be measured, but the basic limitation can not be overcome. We will compare our approach to theirs in more detail in a future paper. \citet{rajbokall01} also considered a number of specific cases of confusion on SIM astrometry. These authors introduced a graphical analogy using phasors as an aid to understanding how errors in the target position arise from confusing sources in the FOV. Typical target fields were constructed on a simulated image with grid spacing of 5 mas, and the amplitude and phase of the fringe which would be measured with SIM for a given wavelength on that image field was computed with a Fourier transform. Diffraction effects at the edges of the (presumed $\approx 1''$ square) SIM FOV were included in constructing the model image, and vector averaging of the individual (narrow) SIM wavelength channels was used to simulate the 1-dimensional apodization of the fringes over the FOV caused by the decreasing coherence of the fringes as the bandwidth increases. \citet{rajbokall01} were particularly interested in modeling the effects of mispointing of the FOV in subsequent visits to the same target field when target proper motions were being measured; in that case the actual distribution of field stars changes as some disappear from one side of the FOV and others appear at the other side. They included photon noise, and also addressed the issue of how the size of the FOV defined by the field stop influenced the level of confusion. Their source field models were constructed to simulate SIM observations of the position and proper motion of target stars in M31, the LMC, and the Galactic bulge. They concluded that the confusion-induced errors in position can often be significant (several times the measurement precision) for faint target stars but the proper motion errors are likely to be small. The errors are, as expected, smaller for wider measurement bands. In one other confusion-related study, \citet{takvellin05} considered the effect of circumstellar disks on the measurement of stellar wobble during observations aimed at detecting extra-solar planets. Their models showed that neither the motion of the disk mass center nor the contamination by disk light is a serious threat to detecting planets around pre-main sequence stars; the basic reason for the insensitivity of these observations to confusion from circumstellar disks is that interferometers tend to resolve such an extended source, reducing its influence on the astrometry of the parent star. The studies summarized above have shown that confusion poses limits to the accuracy of any single SIM measurement, and have therefore succeeded in raising our awareness of this problem. However, the detailed design of SIM has changed since those studies were done, and many of the changes will affect on the modeling results. The size of the collector and its central obscuration, the entrance aperture field stop defining the geometrical FOV, the transmission efficiency of the optics, the fringe disperser design (which defines the bandwidth and central wavelength of the spectral channels), and the QE and spectral response of the detector are now all much better defined. Furthermore, previous studies have focused on specific science programs which were already suspected to be pushing the capabilities of the instrument; they have not provided us with any general ``tools'' for understanding and recognizing confusion, or for dealing with it. Previous studies have also often taken a statistical approach which is less suitable for answering direct questions about specific fields, such as: is a SIM observation of this particular target embedded in that particular field of stars likely to be confused? And, can the observation be done in such a way so as to reduce the confusion bias? What \textit{a priori} information about the target and the field would help? And, if the observation has already been taken, can we identify the effects of confusion in the data? These questions have provided the motivation for the work described in this paper. This paper is organized as follows. In the next section, we summarize the basic Michelson interferometer response as it applies to SIM. We then recall the phasor model introduced by \citet{rajbokall01} and elaborate upon it as a tool for understanding the behavior of confusion in SIM astrometry. Using this analytic model, together with updated knowledge of SIM's instrument parameters, we have constructed a simulation code for evaluating the likelihood of confusion bias in any specific field; details are presented in Section~\ref{simu}. In Section~\ref{limit_values}, we present single measurement confusion bias as a function of magnitude difference and projected separation of an additional star present within the SIM FOV. In Section~\ref{applications}, we apply this semi-analytic model to a number of target fields drawn from the Key Projects which have already been chosen for the initial SIM science program. From this experience we then consider how the single-measurement confusion bias might be reduced through the addition of other information. The most useful additions appear to be knowledge of the approximate locations and spectral energy distributions (SEDs) of the target and of the most troublesome confusing stars in the SIM FOV. Finally, in Section \ref{limitations}, we describe the limitations of our current approach. These limitations are primarily related to the simplified model of the focal plane of SIM which we have used here. In particular, in this paper, we have not modeled the detailed mechanism by which the spectral dispersion is implemented,\footnote{A thin prism disperser turns SIM into an objective prism spectrograph on the CCD detector.} nor have we considered the pixellation of the focal plane by the CCD detector. We have explored these points with the aid of a more elaborate model that includes these effects, and find that the biases estimated using this more elaborate model differ only by small amounts from those provided by the approach described here.\footnote {However, consideration of this more sophisticated model does suggest additional ways to reduce any confusion bias.} We have therefore chosen to present the main issues relevant to SIM confusion with a minimum of complication, and leave the discussion of the more elaborate instrument model to a future paper. Binary stars will be an important class of targets for SIM, and are the topic of one of the major Key Projects. In these cases, the two stars are in a bound orbit and are physically close to each other. Typical binaries to be studied with SIM will have separations from about a few mas to 1000 mas and orbital periods from a few days to several years. Stars in crowded fields can occasionally mimic the effects of binaries if their projected separations become small for particular baseline orientations, but the effects of confusion from such ``apparent'' binaries can be reduced (or even eliminated) by rotating the interferometer baseline and repeating the observation. However, for ``real'' binaries, rotation of the baseline is an integral part of the measurement process. The goal of the binary observation is to obtain the characteristics of the orbit, and this is done by measuring the positions of the components for a number of baseline orientations. These targets are sufficiently specialized that we have removed them from the list of crowded-field problems treated in this paper. Such observations treat binaries as ``signal", whereas here we treat them as ``noise". A discussion of astrometry on binary stars with SIM is planned for a future publication. | \label{sum} We have examined the bias that can occur in a single measurement with SIM owing to the presence of field stars within the FOV. In order to accomplish this task we have presented a model for the SIM interferometer, and a description of how SIM carries out a single measurement of the position of an isolated target star. The measured instrument response is then perturbed by adding a field star to the model FOV; the difference in the angles measured in the two cases is called the ``confusion bias''. The extremes of this bias are calculated for the specific (but common) case of a binary system in order to illustrate its main properties. A number of source models are then developed which resemble the fields to be studied by SIM in several of the Key Projects already selected for inclusion in the initial mission science program. An unconfused version of the source model consisting only of the main target star is used as a reference measurement, and the results compared with a measurement made on the fully-populated field. The difference is the confusion bias in a single SIM measurement. Observations are simulated at various orientations of the interferometer baseline, and variants of the full field are examined in order to understand the sensitivity of the bias to structural details in the field. The magnitude of the confusion bias is found to depend on a number of factors, some obvious, others perhaps less so: \begin{itemize} \item the relative brightnesses of the target and the field stars; \item the shapes of the SEDs of the target and field stars; \item the angular separation of the stars from the center of the FOV; \item the angular separation of the field stars from the target star as projected on the interferometer baseline; and, \item the baseline orientation. \end{itemize} The largest contributions to the confusion bias in a crowded field come from a small number of stars having small projected angular separations from the target, but these stars may actually be located outside of the FOV. Field stars which are less than 4 mag fainter than the target and which have projected separations within 100 mas of the target are potentially the most troublesome. The results of this study provides the understanding and the tools required to examine the likelihood of confusion bias in any single measurement with SIM. Unfortunately, data on the field stars in any specific FOV (especially their positions) is not likely to be available with sufficient accuracy to actually remove this bias.\footnote{There are some possible exceptions one could imagine, but we have not explored them further here.} Our study nevertheless suggests some strategies for recognizing the presence of confusion bias and for dealing with it, both in the observation planning stage and in the data reduction stage. These strategies might include the following: \begin{itemize} \item While dealing with crowded fields, avoid fields with star densities in excess of 0.4 stars per square arcsec. \item If avoidance is impossible, evaluate the likelihood of confusion in the field by using the tools developed here. \item If confusion is likely, try to reduce your sensitivity to it by planning the observing program so that data is taken at the least sensitive orientations of the interferometer baseline. \item If too little is known about the specific field, plan to distribute the available observing time over a number of orientations of the interferometer baseline which differ by a few degrees from each other. Inconsistent values in the data set can then be rejected with motivation. \item If confusion is suspected in a given set of observations for which no prior data exists, acquiring new imagery from e.g., speckle or adaptive optics imaging would be useful for building a model. \end{itemize} There is one additional strategy suggested by our more accurate model of the SIM focal plane. The CCD detector in the focal plane of SIM's camera is planned to have pixels which are smaller than the diameter of the FOV. If the data in the individual pixels can be made available, it would be possible to choose e.g.\ only the central pixel, thereby effectively reducing the FOV and possibly attenuating an offending field star. The penalty of fewer target photons could then be offset by a reduction in the level of confusion bias. This possibility will be discussed in more detail in a future paper \citep{sriron07}. We wish to emphasize that the results of this paper refer to a bias which may be present in a single measurement of angular position with SIM. The determination of the full set of astrometric parameters (position, parallax, proper motion) on any SIM target will be done with a number of measurements, reducing the effects of any single anomalous point. Furthermore, the ultimate accuracy of the results depends on an extensive calibration program to determine the instrumental parameters, including the baseline length and orientations for each field observed. | 7 | 10 | 0710.4262 |
0710 | 0710.4054_arXiv.txt | Boxy/peanut bulges in disk galaxies have been associated to stellar bars. In this talk, we discuss the different properties of such bulges and their relation with the corresponding bar, using a very large sample of a few hundred numerical N-body simulations. We present and inter-compare various methods of measuring the boxy/peanut bulge properties, namely its strength, shape and possible asymmetry. Some of these methods can be applied to both simulations and observations. Our final goal is to get correlations that will allow us to obtain information on the boxy/peanut bulge for a galaxy viewed face-on as well as information on the bars of galaxies viewed edge-on. | Simulations have shown that bars are not vertically thin morphological features, but have a considerable vertical extent and a vertical structure, known as the Boxy/Peanut bulges (hereafter B/P; Combes \& Sanders 1981, Combes et al. 1990). Comparisons between observations and $N$-body simulations have established this direct connection firmer (Athanassoula 2005 and references therein). Furthermore, observations have shown that both bars and B/P bulges are quite predominant in disc galaxies and that the corresponding frequencies are in good agreement with the link between the two structures (L\"utticke, Dettmar \& Pohlen 2000). We measure the peanut properties in a large sample of several hundred $N$-body simulations ran by one of us (EA) for different purposes. More information on these simulations and on their properties can be found in Athanassoula \& Misiriotis (2002) and Athanassoula (2003, 2007). In particular, we seek correlations between the properties of the bar and the properties of the B/P bulge. | We presented several methods to calculate the strength of the bar and the strength, shape and asymmetry of the B/P bulge and found strong correlations between their results. The most important correlation relates the strength of the bar with the strength of the B/P bulge, the strongest bars having the strongest peanuts. We also find that the strength of the peanut depends on the number of buckling episodes it underwent, the strongest bars having undergone more buckling episodes (Fig. 4). Finally, we find a very interesting result about $C_{2,z}(R)$, i.e. about the shape of the radial density profiles along cuts perpendicular to the equatorial plane. For strong bars, having a strong peanut or X-shaped bulge, this profile is more flat-topped, while for weaker bars, with more boxy-like bulges, it is more peaked. All the results summarised here are discussed in length by Athanassoula \& Martinez-Valpuesta (2007, in preparation). \begin{figure} \begin{center} \includegraphics[scale=0.35]{Martinez-Valpuesta_fig4.eps} \caption{Correlations between bar and peanut properties. Each symbol corresponds to one simulation. The type of symbol is related to the number of buckling events suffered by the bar during its evolution. {\it Left panel}: Strength of the B/P bulge measured with our Fourier based method vs. the strength of the bar. {\it Right panel}: Shape of the B/P bulge (i.e. shape of the radial density profiles along cuts perpendicular to the equatorial plane, measured by the minimum of the kurtosis) plotted as a function of bar strength. } \end{center} \end{figure} | 7 | 10 | 0710.4054 |
0710 | 0710.0441_arXiv.txt | We study the non-thermal emissions in a solar flare occurring on 2003 May 29 by using RHESSI hard X-ray (HXR) and Nobeyama microwave observations. This flare shows several typical behaviors of the HXR and microwave emissions: time delay of microwave peaks relative to HXR peaks, loop-top microwave and footpoint HXR sources, and a harder electron energy distribution inferred from the microwave spectrum than from the HXR spectrum. In addition, we found that the time profile of the spectral index of the higher-energy ($\gsim 100$ keV) HXRs is similar to that of the microwaves, and is delayed from that of the lower-energy ($\lsim 100$ keV) HXRs. We interpret these observations in terms of an electron transport model called {\TPP}. We numerically solved the spatially-homogeneous {\FP} equation to determine electron evolution in energy and pitch-angle space. By comparing the behaviors of the HXR and microwave emissions predicted by the model with the observations, we discuss the pitch-angle distribution of the electrons injected into the flare site. We found that the observed spectral variations can qualitatively be explained if the injected electrons have a pitch-angle distribution concentrated perpendicular to the magnetic field lines rather than isotropic distribution. % | \label{sec1} Observations of hard X-rays (HXRs), microwaves, and occasionally gamma-rays in solar flares tell us that a significant amount of non-thermal particles are produced. Among them, HXR and microwave observations are believed to provide the most direct information on electrons. Because HXRs below $\sim 100$ keV are emitted primarily by electrons with energy below several hundred keV via {\bremss} radiation \citep{1971SoPh...18..489B}, whereas microwaves above $\sim 10$ GHz are emitted by electrons above several hundred keV via {\gyros} \citep{1969ApJ...158..753R,1999spro.proc..211B}, these two sources of emission give us information on electrons in two different energy ranges. Therefore, a comparative study by using both HXR and microwave observations is useful for discussing the physics of flare non-thermal electrons over a wide range of energies. Impulsive behavior is commonly seen in both HXR and microwave lightcurves \citep{1974IAUS...57..105K}, but the two emissions do not necessarily behave identically. Temporally, higher-energy HXR and microwave emissions tend to be delayed from lower-energy HXRs \citep[e.g.,][]{1978ApJ...223..620C,1983Natur.305..292N,1985ApJ...292..699B,1997ApJ...487..936A}. \cite{1997ApJ...487..936A} statistically analyzed the low-pass filtered HXR lightcurves for 78 flares observed with the {\it Compton Gamma Ray Observatory} ({\it CGRO}) and find a systematic increase of time delay toward higher energy. They interpreted these time delays in terms of electron precipitation under Coulomb collisions. Spatially, microwave sources do not always coincide with HXR sources. HXRs are typically emitted at the footpoint regions of the flare loop \citep{1994PhDT.......335S} whereas microwaves are emitted mainly at the loop-top region \citep{2002ApJ...580L.185M}. \cite{2002ApJ...580L.185M} suggested that only electrons with a pancake pitch-angle distribution concentrated transverse to the magnetic field lines can explain the observed loop-top microwave source. Spectrally, \cite{2000ApJ...545.1116S} statistically studied the correlation of the HXR and microwave spectral indices for 57 peaks of the non-thermal emission in 27 flares. They found that the electron energy distribution inferred from the microwave spectrum is systematically harder than that inferred from the HXR spectrum, and suggested that the electron energy distribution becomes harder towards higher energy. There are three probable explanations for such spectra: (1) two (or more) different electron populations with distinct physical characteristics, (2) ``second-step acceleration'' \citep[e.g.,][]{1976SoPh...49..343B}, and (3) ``{\TPP} (TPP)'' \citep[e.g.,][]{1976MNRAS.176...15M}. \cite{1976MNRAS.176...15M} presented analytic solutions of the electron energy continuity equation under two conditions: {\it strong} and {\it weak diffusion limits} \citep{1966JGR....71....1K}. In the strong diffusion limit, electrons injected into the loop undergo significant scattering and then are quickly isotropized during the loop transit. They can escape from the loop with a precipitation rate proportional to their velocity, $\nu_{\rm p} \propto v$. In the weak diffusion limit, on the other hand, electrons are less scattered during the transit. When the loss cone distribution is formed, the pitch-angle diffusion time $\tau_{\rm d}$, which is longer than the transit time, controls the electron precipitation, yielding $\nu_{\rm p} \propto 1/\tau_{\rm d}$. The precipitation rate and the evolution of electrons vary, depending on which condition applies. There have been many observations that can be explained in terms of the TPP model \citep[e.g.,][]{2000ApJ...531.1109L}. \cite{2002ApJ...576L..87Y} reported the Nobeyama Radioheliograph (NoRH; \citealt{1994PROCIEEE...82..705}) observation of a flare occurring on 1999 August 28. The NoRH observation showed clear flare-loop structure and propagating features along the loop. They showed that the microwave spectrum in the optically-thin regime is hard (with spectral index $\sim 1.5$) around the loop-top, and then becomes softer (spectral index $\sim 3.5$) toward the footpoints. Their observation indicates that the higher-energy electrons are efficiently trapped within the loop, supporting the TPP model. For this event, however, no HXR observation was available for comparison with the microwave observation. \cite{2000ApJ...545.1116S} pointed out that the discrepancy of the energy distribution between the HXR and microwave emitting electrons found in their study could be explained by the TPP model. In their study, there was no imaging observation to confirm their suggestion. If the HXR and microwave sources do not coincide spatially, the discrepancy of the energy distribution between the HXR and microwave emitting electrons can be explained by the different spatial distribution between the HXR and microwave emitting electrons as a result of the TPP model. Imaging as well as spectral data at both HXR and microwave wavelengths are essential to confirm the role of TPP on the parent electrons In this paper we analyze the non-thermal emissions of the 2003 May 29 flare by using the {\it Reuven Ramaty High Energy Solar Spectroscopic Imager} \citep[RHESSI;][]{2002SoPh..210....3L}, the Nobeyama Radio Polarimeters \citep[NoRP,][and references theirin]{1985PASJ...37..163N} and NoRH. RHESSI has superior spectroscopic ability from $\sim 3$ keV to $\sim 17$ MeV, providing the HXR spectrum from $\sim 3$ keV to $\sim 300$ keV with a spectral resolution of $\sim 1$ keV and arbitrary energy bands. In previous studies, the temporal evolution of the (HXR) spectrum has been considered in less detail, probably due to instrumental limitations. However, the temporally-resolved analysis of the spectrum is important because non-thermal emissions and thus non-thermal electrons are the most ``time-varying'' objects in solar flares. RHESSI enables us to analyze an accurate, temporally-resolved HXR spectrum above $\sim 100$ keV. Because the HXRs above $\sim 100$ keV are mainly emitted by electrons above $\sim 200$ keV \citep{1996ApJ...464..974A}, RHESSI's well-resolved spectral data below $\sim 300$ keV provides us more accurate information on electrons from tens to hundreds of keV than before. Combining the RHESSI HXR and NoRH/NoRP microwave spectral data allows us to fully cover the electrons from tens to thousands of keV. For a physical interpretation of the observations, we use a numerical model of TPP which treats the pitch-angle diffusion more generally than the analytic solutions developed for the weak and strong diffusion limits. \cite{2000ApJ...543..457L} performed a similar treatment of the electron transport to explain their microwave observation of a flare on 1993 June 3. We also predict the microwave and HXR emissions from the calculated electron distribution. Comparing these model results with the observations, we discuss electron injection and transport, and address how the pitch-angle distribution of the injected electrons affects the evolution of the trapped and precipitating electrons, and their resultant emissions. The paper proceeds as follows. In {\S}~\ref{sec2} we present a comparative study of the non-thermal emissions of a solar flare occurring on 2003 May 29, by using the RHESSI HXR and Nobeyama microwave observations. Temporally-resolved spectra of the HXRs and microwaves are analyzed in detail. We discuss energy-dependent delays of the time profiles of the spectral indices, which have not been discussed in previous studies. In {\S}~\ref{sec3} we present our treatment of the TPP model. We numerically solve the spatially-homogeneous {\FP} equation \citep{1990A&A...230..213M} with the Coulomb interaction \citep[e.g.,][]{1981ApJ...251..781L} and a time-dependent injection. In {\S}~\ref{sec4} we describe the time evolution of the trapped and precipitating electron distribution and the predicted microwave and HXR emissions. The behavior of the HXR and microwave emissions predicted by the model are compared with the observations, allowing us to give some constraints on the properties of flare non-thermal electrons. In {\S}~\ref{sec5} we conclude our study. % | \label{sec5} We presented the comparative study of the non-thermal emissions of the flare occurring on 2003 May 29 using the RHESSI HXR and Nobeyama microwave observations. Further, we considered the electron transport model, TPP, to explain the observations. The 2003 May 29 flare showed two non-thermal HXR sources at the footpoints and a microwave source at the loop-top, as observed with RHESSI and NoRH. We interpreted this in terms of the TPP model. We presented the time profiles of the spectral indices of the higher-energy HXRs $\gammat{H}{obs}$ as well as the lower-energy HXRs $\gammat{L}{obs}$ and microwaves $\alphat{obs}$. The spectra of microwaves and HXRs imply that the microwave emitting electrons have a harder energy distribution than the HXR emitting ones. % We found that the time profile of $\gammat{H}{obs}$ shows similarity with that of $\alphat{obs}$ rather than with $\gammat{L}{obs}$, and is delayed from that of $\gammat{L}{obs}$. We numerically solved the spatially-homogeneous {\FP} equation for the TPP model to describe the evolution of electrons. Precipitating electrons have a softer energy distribution than the trapped ones in the weak diffusion regime. Differences of the injection pitch-angle distribution especially affect the evolution of the precipitating electrons. We calculated the microwave and HXR emissions from the calculated trapped electron distribution and precipitation flux for comparison with the observations. The TPP model in the weak diffusion regime can yield a soft HXR spectrum and a hard microwave spectrum. The calculated difference of the spectral indices between the HXRs and microwaves, $\sim 1.5$, is in agreement with the observations. We further found that a pancake pitch-angle distribution for the injected electrons rather than an isotropic distribution is more adequate to qualitatively explain the temporal variation of $\gammat{H}{obs}$. By comparing the model calculation with the observation, we can constrain the pitch-angle distribution of the injected electrons, which is crucially important for understanding the electron acceleration mechanism in solar flares. Currently, we are improving our treatment of the TPP model to include the spatial inhomogeneity in the Fokker-Planck equation. Using this, a systematic investigation of the best parameter set to explain the observation is in progress, and will be reported in the future. | 7 | 10 | 0710.0441 |
0710 | 0710.0993_arXiv.txt | The Alpha Magnetic Spectrometer (AMS), whose final version AMS-02 is to be installed on the International Space Station (ISS) for at least 3 years, is a detector designed to measure charged cosmic ray spectra with energies up to the TeV region and with high energy photon detection capability up to a few hundred GeV, using state-of-the art particle identification techniques. It is equipped with several subsystems, one of which is a proximity focusing Ring Imaging \CK\ (RICH) detector equipped with a dual radiator (aerogel+NaF), a lateral conical mirror and a detection plane made of 680 photomultipliers and light guides, enabling precise measurements of particle electric charge and velocity ($\Delta \beta / \beta \sim$ 10${}^{-3}$ and 10${}^{-4}$ for $Z=$~1 and $Z=$~10~$-$~20, respectively) at kinetic energies of a few GeV/nucleon. Combining velocity measurements with data on particle rigidity from the AMS-02 Tracker ($\Delta R / R \sim$ 2\% for $R=$~1~$-$~10 GV) it is possible to obtain a reliable measurement for particle mass. One of the main topics of the AMS-02 physics program is the search for indirect signatures of dark matter. Experimental data indicate that dark, non-baryonic matter of unknown composition is much more abundant than baryonic matter, accounting for a large fraction of the energy content of the Universe. Apart from antideuterons produced in cosmic-ray propagation, the annihilation of dark matter will produce additional antideuteron fluxes. Detailed Monte Carlo simulations of AMS-02 have been used to evaluate the detector's performance for mass separation, a key issue for $\bar{D}/\bar{p}$ separation. Results of these studies are presented. | The Alpha Magnetic Spectrometer (AMS)\cite{bib:ams}, whose final version AMS-02 is to be installed on the International Space Station (ISS) for at least 3 years, is a detector designed to study the cosmic ray flux by direct detection of particles above the Earth's atmosphere using state-of-the-art particle identification techniques. AMS-02 is equipped with a superconducting magnet cooled by superfluid helium. The spectrometer is composed of several subdetectors: a Transition Radiation Detector (TRD), a Time-of-Flight (TOF) detector, a Silicon Tracker, Anticoincidence Counters (ACC), a Ring Imaging \CK\ (RICH) detector and an Electromagnetic Calorimeter (ECAL). Fig.~\ref{amsdet} shows a schematic view of the full AMS-02 detector. A preliminary version of the detector, AMS-01, was successfully flown aboard the US space shuttle Discovery in June 1998. \begin{figure}[htb] \center \mbox{\epsfig{file=ams2.eps,width=0.48\textwidth,clip=}} \caption{Exploded view of the AMS-02 detector.\label{amsdet}} \end{figure} The main goals of the AMS-02 experiment are: \begin{itemize} \item A precise measurement of charged cosmic ray spectra in the rigidity region between \mbox{$\sim$ 0.5 GV} and \mbox{$\sim$ 2 TV}, and the detection of photons with energies up to a few hundred GeV; \item A search for heavy antinuclei ($Z \ge$ 2), which if discovered would signal the existence of cosmological antimatter; \item A search for dark matter constituents by examining possible signatures of their presence in the cosmic ray spectrum. \end{itemize} The long exposure time and large acceptance (0.5 m${}^2$sr) of AMS-02 will enable it to collect an unprecedented statistics of more than $10^{10}$ nuclei. \begin{figure}[htb] \center \mbox{\epsfig{file=ev702white.eps,width=0.48\textwidth,clip=}} \vspace{-0.3cm} \caption{A simulated proton event as seen in the AMS-02 display.\label{amsdisplay1}} \end{figure} | AMS-02 will provide a major improvement on the current knowledge of cosmic rays. A total statistics of more than 10${}^{10}$ events will be collected during its operation. Detailed simulations have been performed to evaluate the detector's particle identification capabilities, in particular those of the RICH, which might be crucial for the identification of an antideuteron flux resulting from neutralino annihilation. Simulation results show that the separation of light isotopes is feasible. Using a set of simple cuts based on event data, relative mass resolutions of $\sim$~2 \% and rejection factors up to 10${}^4$ have been attained in D/p separation at energies of a few GeV/nucleon. \newpage | 7 | 10 | 0710.0993 |
0710 | 0710.5578_arXiv.txt | We investigate properties of iron fluorescent line arising as a result of illumination of a black hole accretion disc by an X-ray source located above the disc surface. We study in details the light-bending model of variability of the line, extending previous work on the subject. We indicate bending of photon trajectories to the equatorial plane, which is a distinct property of the Kerr metric, as the most feasible effect underlying reduced variability of the line observed in several objects. A model involving an X-ray source with a varying radial distance, located within a few central gravitational radii around a rapidly rotating black hole, close to the disc surface, may explain both the elongated red wing of the line profile and the complex variability pattern observed in MCG--6-30-15 by {\it XMM-Newton}. We point out also that illumination by radiation which returns to the disc (following the previous reflection) contributes significantly to formation of the line profile in some cases. As a result of this effect, the line profile always has a pronounced blue peak (which is not observed in the deep minimum state in MCG--6-30-15), unless the reflecting material is absent within the innermost 2--3 gravitational radii. | Broad iron lines observed from many black hole systems most likely originate from the innermost regions of an accretion disc and their profiles are shaped by gravitational redshift and Doppler shifts. Modelling of the lines observed in several objects requires strongly enhanced fluorescent emission from a few gravitational radii (e.g., Wilms et al.~2001, Fabian et al.~2002, Miller et al.~2002, Miller et al.~2004; see review in Reynolds \& Nowak 2003), which in turn indicates that a primary source of hard X-ray emission must also be located close to the black hole. Thus, both the primary and reflected emission should be subject to strong gravity effects. These effects are also tentatively considered as an explanation of the complex variability pattern characterising radiation reflected from disc, including the iron line, and the primary hard X-ray continuum emission. Namely, weak variability of the reflected component, uncorrelated with the variability of the primary emission, has been reported in a number of sources (e.g., Vaughan \& Fabian 2004; see review in Fabian \& Miniutti 2005). This is contrary to expectations from a simple geometrical model of a hard X-ray source located close to the reflecting disc, where a strict correlation between variations of the primary and reflected emission should be observed. Fabian \& Vaughan (2003) first argued that such a reduced variability may be explained by relativistic effects, in particular by light bending and focusing the primary emission towards the accretion disc. Qualitatively, variations of the reflected emission should be much weaker as changes of the height of the X-ray source cause variations of its observed luminosity at infinity, while changes of the flux received by the disc (enhanced by the gravitational focusing; e.g., Matt et al.\ 1992; Martocchia \& Matt 1996; Petrucci \& Henri 1997) are much weaker. However, for a static primary source located on the symmetry axis, the illuminating radiation is focused into the innermost part of the disc and the reflected emission is subject to similar light bending as the primary. As a result, similar variability characterises the primary and reflected emission, at least for observers with low inclination angles. Miniutti et al.\ (2003) and Miniutti \& Fabian (2004) have developed further this model to include rotation of the primary source around the axis and the resulting beaming of its emission toward outer regions of the disc. Then, variations of the reflected emission are reduced because of the more extended region it is produced. Predictions of their model are found to be consistent with observations of black-hole systems, e.g.\ by Miniutti et al.\ (2003), Miniutti, Fabian \& Miller (2004), Fabian et al.\ (2004). However, a number of important effects were not systematically studied. In particular, a specific pattern of motion of a primary source is assumed (including both location relative to the symmetry axis and azimuthal motion) but it is not discussed how strongly the resulting properties depend on these assumptions. For example, corotation of the primary source with the disc is assumed by Miniutti \& Fabian (2004), which is likely close to the disc surface, but at most approximate at high latitudes. As the azimuthal motion of the primary source relative to the disc may affect significantly the reflected emission (see e.g.\ Reynolds \& Fabian 1997), it is not clear whether predictions of this model are specific to the underlying assumptions or generic to models involving the light bending. In this paper we systematically analyse the light bending model, concentrating on strong gravity effects. We neglect some other effects which may contribute to the original variability problem, for example, ionization of the disc surface (Nayakshin \& Kazanas 2002). We focus here on the iron K$\alpha$ line; an analysis of reflected emission including the Compton reflected radiation will be presented in our next paper. We study in details a number of geometrical scenarios of the source location and motion. We point out certain inadequacies in the original computations of Miniutti \& Fabian (2004), and how they influence their quantitative results. We find also a novel scenario in which reduction of the line variability follows directly from properties of photon transfer in the Kerr metric. Namely, we find that a source located close to the disc surface, with a varying radial distance from a Kerr black hole, gives rise to both an approximately constant illumination of the surrounding disc and a very strong variability of the primary emission observed at infinity. Some aspects of variability in similar models have been considered recently e.g.\ by Czerny et al.\ (2004) and Pech\'acek et al.\ (2005), however, these studies considered only emission emerging locally from the region under the source. On the other hand, we find that transfer of primary emission to more distant regions of the disc is crucial for variability effects in such a model. We concentrate on low inclination objects, with application to Seyfert 1 galaxies in mind. Obviously, the Doppler shifts are more pronounced at high inclinations, but observational studies of high inclination objects are less advanced, either because they are obscured (as Seyfert 2 galaxies), or because the studies on dynamical time scale are not possible (as in stellar mass black hole systems). Furthermore, we consider only emission averaged over at least an orbital period. Effects resulting from varying azimuthal location of an off-axis source, with respect to observer, are studied, e.g., in Ruszkowski (2000), Yu \& Lu (2000) and Goyder \& Lasenby (2004). In Section 3 we analyse various physical effects relevant to formation of relativistic line profiles and variability effects; in Section 4 we apply our results to a Seyfert 1 galaxy MCG--6-30-15 which is the best studied object with clear signatures of strong gravity effects. | \subsection{Variability models} We have extended the model, formulated by Miniutti \& Fabian (2004), relating reduced variability of the reflected emission to changing magnitude of relativistic effects as location of the primary X-ray source changes. We find that original computations of Miniutti \& Fabian (2004) - for a model involving a vertically moving source - overestimate the reduction effect by assuming a value of the outer radius of the disc which is too small for the range of the source heights considered in that model. On the other hand, we find a significant reduction of the variability of reflected emission in a model with a rapidly rotating black hole and a source moving radially, low above the disc surface. The reduced variability occurs then for the innermost range of radial distances, $\le 4R_{\rm g}$. We find also that - only in this range of parameters - the GR effects give rise to a significant decline around 6.5 keV in rms spectra. Note that generation of strong X-ray emission at these distances ($\le 4R_{\rm g}$) is consistent with the condition of a high value of $a$, as - for a rapidly rotating black hole - most accretion power is dissipated within a few innermost $R_{\rm g}$. \subsection{The black hole spin} Determination of the value of black hole spin or, even more fundamentally, verification of effects predicted by the Kerr metric solution of GR equations - remains a major issue of black hole astrophysics. In this context, several effects are taken into account for X-ray spectroscopy. Strong redshift of photons forming the observed red wings, at $E<4$ keV, is considered as evidence of rapid rotation of a black hole (e.g., Brenneman \& Reynolds 2006). However, similar redshift can be obtained for $a=0$ if emission from $r_{\rm d}<6$ is taken into account (Reynolds \& Begelman 1997). Then, the derived high value of $a$ relies on assumption of no neutral iron emission from within the radius of marginal stability. Response of the line to increase of primary emission has been suggested for future studies. In particular, Reynolds et al.\ (1999) indicate a bump occurring in the line profile and progressing to lower energies, with proceeding time, as a feature characteristic for a rapidly rotating black hole. Again, a similar - redshifted and Shapiro delayed - bump should appear for a non-rotating black hole if fluorescence inside $r_{\rm ms}$ was taken into account. A straightforward analysis of the space-time metric could be performed for systems observed close to edge-on, for which effects due to lensing by a black hole would be directly seen in the line profile (e.g., Zakharov \& Repin 2003). However, as noted in Narayan \& McClintock (2005), there seems to be a selection effect preventing such systems from being observed. Then, we point out that a (largely unambiguous) analysis of imprints of the space-time metric would be possible in profile of the line resulting from irradiation by a strong flare just above the disc surface. If such a flare occurred at $r_{\rm s} \la 4$, properties of the space-time related with the black hole rotation would result in $\ga 10$ per cent in magnitude effects in the line profile. The effects related to the value of $a$ are rather subtle but in principle possible to establish observationally. A compact flare dominating total emission would be required to make such an analysis feasible and viability of such scenario is uncertain. Such flares are indeed occasionally observed (e.g., Ponti et al.\ 2004). Interestingly, however, the Fe K$\alpha$ line was actually very weak during the flare analysed by Ponti et al.\ (2004), and a strong line appeared in the spectrum with significant time delay after the flare. Finally, we emphasise that the reduced variability of reflected component - basing on mechanism advocated in this paper - is itself a direct manifestation of the nature of the Kerr space-time. Note that another effect resulting from properties of the Kerr metric, namely anisotropic emission of hard X-rays generated close to a rotating black hole, is qualitatively consistent with inclination-angle dependence of intrinsic spectra of Seyfert galaxies (Nied\'zwiecki 2005). In support for the tentative relation of these two effects to properties of the Kerr metric, note also that comparison of the total mass in black holes in the local universe with the total luminosity produced by active galactic nuclei indicates that most supermassive black holes should rotate rapidly, e.g.\ Elvis, Risaliti \& Zamorani (2002). \subsection{Modelling the red wing of relativistic Fe lines} As discussed above, detection of photons with $E < 4$ keV is considered as the evidence of rapid rotation of the black hole. Such strongly redshifted photons were revealed in the Fe line profiles observed in MCG--6-30-15 and several other objects (e.g., Miller et al.\ 2004, Miniutti et al.\ 2004). When fitted by models assuming a power-law radial emissivity, the observed profiles require $q>4$ within $(6-10) R_{\rm g}$. Considering physical scenarios for generation of such profiles, we find that they \newline (i) can be produced as a result of illumination by a hard X-ray source located close to a black hole (preferably at $r_{\rm s} =2$--3) and rather close to the disc surface ($h_{\rm s} \la 1$); \newline (ii) cannot be explained by models involving a source at a height of several $R_{\rm g}$, or higher, close to the axis; \newline if they are to come from a disc, with small inclination, surrounding a rotating black hole. Regarding (ii) we find, in agreement with previous studies, that such location of the source yields a steep $\epsilon_{\rm Fe}$, with $q>3$, in the innermost part of the disc. However, this steep $\epsilon_{\rm Fe}$ occurs only within $r_{\rm d} \le 2$ (and due to gravitational blueshift rather than light bending). The majority of emission from this region of the disc is either captured or bent toward the disc plane; moreover, the observed photons are strongly redshifted. As a result, this emission gives only a minor contribution to the line profile and it is not relevant to the shape and strength of the observed red wings. This conclusion is not affected by effects related with azimuthal motion of the source, returning radiation or angular emission law of fluorescence; these effects influence mostly the strength of the blue peak. On the other hand, the red wing is primarily related with position of the X-ray source. \subsection{Returning radiation, non-local irradiation} A further challenge for modelling certain line profiles results from illumination of the surrounding disc (i.e. beyond a few $R_{\rm g}$) by returning radiation. We find that this effect may be strong for the range of parameters (small $r_{\rm s}$ and $V$) previously not explored. A similarly strong enhancement of Compton reflection by returning radiation, for very small distances ($r_{\rm s} \la 2$) of primary source, was noted recently by Suebsewong et al.\ (2006). Another effect, related to bending of photon trajectories to the disc plane, affects models relating the Fe radial profile to some physical processes. These models typically assume that the radial profile of Fe emission is the same as some physically motivated profile of energy dissipation. E.g., Reynolds et al.\ (2004) make such assumption in their analysis of the {\it XMM} observation of MCG--6-30-15, applying model of the torqued-disc emission (where emission results mostly from extraction of rotational energy of the black hole and thus is very centrally concentrated). Moreover, they assume that there is no dissipation, and thus no reprocessing, beyond rather small (several $R_{\rm g}$) outer radius. We emphasise that even for X-rays generated very low above the disc surface, a significant illumination of more distant regions of the disc should occur. Both the returning radiation and the non-locality of irradiation result primarily in enhancement of the blue peak of the line. \subsection{Azimuthal motion} Impact of strong gravity is usually studied under specific assumptions on the X-ray source motion. Typically, angular velocity of the source is related to that of the disc (e.g. Ruszkowski 2000; Miniutti \& Fabian 2004), while Dabrowski \& Lasenby (2001) assume a static primary source. We note that change of $V$ significantly affects flux and profile of the iron line through combination of SR and GR effects. Moreover, the Kerr metric terms yield a non-trivial dependence between $V$ and $\Omega$ (equation (\ref{v})). Then, the assumed parametrisation of motion may be crucial for derived properties, e.g.\ in some variability models. Obviously, additional effects may result from relativistic vertical or radial motion of the source (e.g.\ Yu \& Lu 2001), which were neglected in this paper. \subsection{Applicability of our results} We focused above on MCG--6-30-15, for which several observed properties may be explained by our model. Similar effects, including reduced variability of reflection, pronounced red wings or strongly reflection dominated spectra, have been revealed in some Narrow Line Seyfert 1 galaxies (e.g. Fabian et al.\ 2004, Ponti et al.\ 2006), as well as in stellar mass black-hole systems (e.g.\ Miniutti et al.\ 2004) and in galactic nuclei with much higher black hole masses (Fabian et al.\ 2005). Note that our scenario implies that the X-ray luminosity is underestimated by an order of magnitude. in these objects. On the other hand, most black-hole systems do not show signatures of such extreme effects. Then, similar reduction of the X-ray flux does not necessarily characterise any low inclination system. Note, however that such property would be inevitable for rapidly rotating black holes, where major fraction of accretion power is dissipated at very small $r_{\rm s}$, if this power is converted into hard X-rays in situ. As in virtually all previous studies of that subject, we analysed impact of GR effects for the simplest case of an isotropic point source of hard X-ray emission, which approach follows from our lack of understanding of the X-ray source nature. The derived properties result from transfer of radiation from source to disc and observer and from disc to observer. Additional (but smaller in magnitude) effects may occur in more realistic models. E.g., Comptonization close to a Kerr black hole gives rise to anisotropic emission, cf.~Nied\'zwiecki (2005), therefore radiation with different spectral index may irradiate various parts of the disc, while the transfer effects affect only normalisation and cut-off energy of the primary emission. The simplified description of X-ray source considered in this paper approximates most closely scenario with magnetic flares above the disc surface. Then, our results may be directly applicable to a model with flares occurring randomly at various radial distances. Qualitatively, we may assess similar variability effects for continuous spacial distributions of the hard X-ray source, e.g.\ an extended corona covering the disc surface or a small hot torus replacing the disc within a few innermost $R_{\rm g}$. Varying size of such an extended, hot plasma, in the Kerr metric, should give rise reduced variability of reflected component. In particular, a decline of direct emission from a shrinking corona or torus would be observed, even if its intrinsic luminosity remained unchanged, while increasing fraction of its emission would be bent to the disc plane giving rise to strong reflection component. The major shortcoming in our study results from the neglect of ionization effects. The strongest ionization should occur below the source, especially in models with low height above the disc surface. In general, strong ionization of this region should suppress the redshifted Doppler horns formed in the red-wing. On the one hand, this would make studies of effects related to value of $a$ less feasible. On the other hand, such depletion of the variable contribution to the red wing would reduce variations of the line profile at various flux states. However, details of ionization structure, and of the related reduction of the variable contribution to the lien, depend on additional assumptions, in particular, on azimuthal distribution of flares. Obviously, a strong single flare would give rise to stronger ionization than many weak flares uniformly distributed at a given $r_{\rm s}$. \subsection{Summary} If attributed to strong-gravity effects, both the time-averaged line profile and the variability pattern observed in MCG--6-30-15 by {\it XMM-Newton\/} independently indicate that a primary hard X-ray source must be located very close ($\la 4R_{\rm g}$) to a black hole, i.e.\ in the region where Kerr metric effects become crucial. Rapid rotation of the black hole is necessary (and vastly sufficient) to account for reduction of variability of reflected emission for such spacial location of the source. Bending to the equatorial plane, underlying this reduction effect, appears to be the most pronounced effect of the Kerr metric to be studied in the X-ray spectra of black-hole systems. | 7 | 10 | 0710.5578 |
0710 | 0710.5052_arXiv.txt | Magnetic fluctuations generated by a tangling of the mean magnetic field by velocity fluctuations are studied in a developed turbulent convection with large magnetic Reynolds numbers. We show that the energy of magnetic fluctuations depends on magnetic Reynolds number only when the mean magnetic field is smaller than $B_{\rm eq} / 4 {\rm Rm}^{1/4}$, where $B_{\rm eq}$ is the equipartition mean magnetic field determined by the turbulent kinetic energy and ${\rm Rm}$ is magnetic Reynolds number. Generation of magnetic fluctuations in a turbulent convection with a nonzero mean magnetic field results in a decrease of the total turbulent pressure and may cause formation of the large-scale inhomogeneous magnetic structures even in an originally uniform mean magnetic field. This effect is caused by a negative contribution of the turbulent convection to the effective mean Lorentz force. The inhomogeneous large-scale magnetic fields are formed due to the excitation of the large-scale instability. The energy for this instability is supplied by the small-scale turbulent convection. The discussed effects might be useful for understanding the origin of the solar nonuniform magnetic fields, e.g., sunspots. | Magnetic fields in astrophysics are strongly nonuniform (see, e.g., \cite{M78,P79,KR80,ZRS83,RSS88,RH04,O03,BS05}). Large-scale magnetic structures are observed in the form of sunspots, solar coronal magnetic loops, etc. There are different mechanisms for the formation of the large-scale magnetic structures, e.g., the magnetic buoyancy instability of stratified continuous magnetic field \cite{P79,P66,G70,P82}, the magnetic flux expulsion \cite{W66}, the topological magnetic pumping \cite{DY74}, etc. Magnetic buoyancy applies in the literature for different situations (see \cite{P82}). The first corresponds to the magnetic buoyancy instability of stratified continuous magnetic field (see, e.g., \cite{P79,P66,G70,P82}), and magnetic flux tube concept is not used there. The magnetic buoyancy instability of stratified continuous magnetic field is excited when the scale of variations of the initial magnetic field is less than the density stratification length. On the other hand, buoyancy of discrete magnetic flux tubes has been discussed in a number of studies in solar physics and astrophysics (see, e.g., \cite{P82,P55,S81,SB82,FS93,SC94}). This phenomenon is also related to the problem of the storage of magnetic fields in the overshoot layer near the bottom of the solar convective zone (see, e.g., \cite{SW80,T01,TH04,B05}). A universal mechanism of the formation of the nonuniform distribution of magnetic flux is associated with a magnetic flux expulsion. In particular, the expulsion of magnetic flux from two-dimensional flows (a single vortex and a grid of vortices) was demonstrated in \cite{W66}. In the context of solar and stellar convection, the topological asymmetry of stationary thermal convection plays very important role in the magnetic field dynamics. In particular, the topological magnetic pumping is caused by the topological asymmetry of the thermal convection \cite{DY74}. The fluid rises at the centers of the convective cells and falls at their peripheries. The ascending fluid elements (contrary to the descending fluid elements) are disconnected from one another. This causes a topological magnetic pumping effect allowing downward transport of the mean horizontal magnetic field to the bottom of a cell but impeding its upward return \cite{ZRS83,DY74,GP83}. Turbulence may form inhomogeneous large-scale magnetic fields due to turbulent diamagnetic and paramagnetic effects (see, e.g., \cite{KR80,Z57,VK83,K91,KR92,RKR03}). Inhomogeneous velocity fluctuations lead to a transport of mean magnetic flux from regions with high intensity of the velocity fluctuations. Inhomogeneous magnetic fluctuations due to the small-scale dynamo cause turbulent paramagnetic velocity, i.e., the magnetic flux is pushed into regions with high intensity of the magnetic fluctuations. Another effects are the effective drift velocities of the mean magnetic field caused by inhomogeneities of the fluid density \cite{K91,KR92} and pressure \cite{KP93}. In a nonlinear stage of the magnetic field evolution, inhomogeneities of the mean magnetic field contribute to the diamagnetic or paramagnetic drift velocities depending on the level of magnetic fluctuations due to the small-scale dynamo and level of the mean magnetic field \cite{RK04}. The diamagnetic velocity causes a drift of the magnetic field components from the regions with a high intensity of the mean magnetic field. The nonlinear drift velocities of the mean magnetic field in a turbulent convection have been determined in \cite{RK06}. This study demonstrates that the nonlinear drift velocities are caused by the three kinds of the inhomogeneities, i.e., inhomogeneous turbulence; the nonuniform fluid density and the nonuniform turbulent heat flux. The nonlinear drift velocities of the mean magnetic field cause the small-scale magnetic buoyancy and magnetic pumping effects in the turbulent convection. These phenomena are different from the large-scale magnetic buoyancy and magnetic pumping effects which are due to the effect of the mean magnetic field on the large-scale density stratified fluid flow. The small-scale magnetic buoyancy and magnetic pumping can be stronger than these large-scale effects when the mean magnetic field is smaller than the equipartition field determined by the turbulent kinetic energy \cite{RK06}. The pumping of magnetic flux in three-dimensional compressible magnetoconvection has been studied in direct numerical simulations in \cite{OS02} by calculating the turbulent diamagnetic and paramagnetic velocities. Turbulence may affect also the Lorentz force of the large-scale magnetic field (see \cite{KRR89,KRR90,KR94,KMR96}). This effect can also form inhomogeneous magnetic structures. In this study a theoretical approach proposed in \cite{KRR89,KRR90,KR94,KMR96} for a nonconvective turbulence is further developed and applied to investigate the modification of the large-scale magnetic force by turbulent convection and to elucidate a mechanism of formation of inhomogeneous magnetic structures. This paper is organized as follows. In Sect.~II we discuss the physics of the effect of turbulence on the large-scale Lorentz force. In Sect.~III we formulate the governing equations, the assumptions, the procedure of the derivations of the large-scale effective magnetic force in turbulent convection. In Sect.~IV we study magnetic fluctuations and determine the modification of the large-scale effective Lorentz force by the turbulent convection. In Sect.~V we discuss formation of the large-scale magnetic inhomogeneous structures in the turbulent convection due to excitation of the large-scale instability. Finally, we draw conclusions in Sect.~VI. In Appendix~A we perform the derivation of the large-scale effective Lorentz force in the turbulent convection. | In the present study we investigate magnetic fluctuations generated by a tangling of the mean magnetic field in a developed turbulent convection. When the mean magnetic field $B \ll B_{\rm eq} / 4 {\rm Rm}^{1/4}$, the energy of magnetic fluctuations depends on magnetic Reynolds number. We study the modification of the large-scale magnetic force by turbulent convection. We show that the generation of magnetic fluctuations in a turbulent convection results in a decrease of the total turbulent pressure and may cause formation of the large-scale magnetic structures even in an originally uniform mean magnetic field. This phenomenon is due to a negative contribution of the turbulent convection to the effective mean magnetic force. The large-scale instability causes the formation of inhomogeneous magnetic structures. The energy for these processes is supplied by the small-scale turbulent convection, and this effect can develop even in an initially uniform magnetic field. In contrast, the Parker's magnetic buoyancy instability is excited when the density stratification scale is larger than the characteristic scale of the mean magnetic field variations (see \cite{P66,P79,G70}). The free energy in the Parker's magnetic buoyancy instability is drawn from the gravitational field. The characteristic time of the large-scale instability is of the order of the Alfv\'{e}n time based on the large-scale magnetic field. We study an initial stage of formation of the large-scale magnetic structures for horizontal and vertical mean magnetic fields relative to the vertical direction of the gravity field. In the turbulent convection there are two ranges for the large-scale instability of the horizontal mean magnetic field. The first range for the instability is related to the negative contribution of turbulence to the effective magnetic pressure for the case of $L_{\rho} < L_{B}$, while the second range for the instability is mainly caused by the anisotropic contribution of the turbulent convection to the effective magnetic force. The large-scale instability of the vertical uniform mean magnetic field is caused by the modification of the mean magnetic tension by small-scale turbulent convection. The discussed effects in the present study might be useful for the understanding of the origin of the sunspot formation. Since in the present study we neglect very small Brunt-V\"{a}is\"{a}l\"{a} frequency based on the gradient of the mean entropy, we do not investigate the large-scale dynamics of the mean entropy. This problem was addressed in \cite{KM00} whereby the modification of the mean magnetic force by the turbulent convection was not taken into account. In order to study magnetic fluctuations and the modification of the large-scale Lorentz force by turbulent convection we apply the spectral $\tau$ approximation (see Sect.~III). The $\tau$ approach is an universal tool in turbulent transport that allows to obtain closed results and compare them with the results of laboratory experiments, observations and numerical simulations. The $\tau$ approximation reproduces many well-known phenomena found by other methods in turbulent transport of particles and magnetic fields, in turbulent convection and stably stratified turbulent flows (see below). In turbulent transport, the $\tau$ approximation yields correct formulae for turbulent diffusion, turbulent thermal diffusion and turbulent barodiffusion (see, e.g., \cite{EKR96,BF03}). The phenomenon of turbulent thermal diffusion (a nondiffusive streaming of particles in the direction of the mean heat flux), has been predicted using the stochastic calculus (the path integral approach) and the $\tau$ approximation. This phenomenon has been already detected in laboratory experiments in oscillating grids turbulence \cite{EEKR04} and in a multi-fan turbulence generator \cite{EEKR06} in stably and unstably stratified fluid flows. The experimental results obtained in \cite{EEKR04,EEKR04} are in a good agreement with the theoretical studies performed by means of different approaches (see \cite{EKR96,PM02}). The $\tau$ approximation reproduces the well-known $k^{-7/3}$-spectrum of anisotropic velocity fluctuations in a sheared turbulence (see \cite{EKRZ02}). This spectrum was found previously in analytical, numerical, laboratory studies and was observed in the atmospheric turbulence (see, e.g., \cite{L67}). In the turbulent boundary layer problems, the $\tau$-approximation yields correct expressions for turbulent viscosity, turbulent thermal conductivity and the classical heat flux. This approach also describes the counter wind heat flux and the Deardorff's heat flux in convective boundary layers (see \cite{EKRZ02}). These phenomena have been studied previously using different approaches (see, e.g., \cite{MY75,Mc90,Z91}). The theory of turbulent convection \cite{EKRZ02} based on the $\tau$-approximation explains the recently discovered hysteresis phenomenon in laboratory turbulent convection \cite{EEKRM06}. The results obtained using the $\tau$-approximation allow also to explain the most pronounced features of typical semi-organized coherent structures observed in the atmospheric convective boundary layers ("cloud cells" and "cloud streets") \cite{ET85}. The theory \cite{EKRZ02} based on the $\tau$-approximation predicts realistic values of the following parameters: the aspect ratios of structures, the ratios of the minimum size of the semi-organized structures to the maximum scale of turbulent motions and the characteristic lifetime of the semi-organized structures. The theory \cite{EKRZ02} also predicts excitation of convective-shear waves propagating perpendicular to the convective rolls ("cloud streets"). This waves have been observed in the atmospheric convective boundary layers with cloud streets \cite{ET85}. A theory \cite{ZEKR07} for stably stratified atmospheric turbulent flows based on the $\tau$-approximation and the budget equations for the key second moments, turbulent kinetic and potential energies and vertical turbulent fluxes of momentum and buoyancy, is in a good agrement with data from atmospheric and laboratory experiments, direct numerical simulations and large-eddy simulations (see detailed comparison in Sect. 5 of \cite{ZEKR07}). The detailed verification of the $\tau$ approximation in the direct numerical simulations of turbulent transport of passive scalar has been recently performed in \cite{BK04}. In particular, the results on turbulent transport of passive scalar obtained using direct numerical simulations of homogeneous isotropic turbulence have been compared with that obtained using a closure model based on the $\tau$ approximation. The numerical and analytical results are in a good agreement. In magnetohydrodynamics, the $\tau$ approximation reproduces many well-known phenomena found by different methods, e.g., the $\tau$ approximation yields correct formulae for the $\alpha$-effect \cite{KR80,RK93,RK00,RKR03}, the turbulent diamagnetic and paramagnetic velocities \cite{Z57,VK83,K91,KR92,RKR03}, the turbulent magnetic diffusion \cite{KR80,VK83,KRP94,RKR03,RK04}, the ${\bf \Omega} {\bf \times} {\bf J}$ effect and the $\kappa$-effect \cite{KR80,RKR03}, the shear-current effect \cite{RK03,RK04,RK07}. Generation of the large-scale magnetic field in a nonhelical turbulence with an imposed mean velocity shear has been recently investigated in \cite{BH05} using direct numerical simulations. The results of these numerical simulations are in a good agreement with the theoretical predictions based on the $\tau$ approximation (see \cite{RK03,RK04,RK07}) and with the numerical solutions of the nonlinear dynamo equations performed in \cite{BS05B,RKL06} (see detailed comparison in \cite{RK07}). The validity of the $\tau$ approximation has been tested in the context of dynamo theory, in direct numerical simulations in \cite{BSM05}. The alpha effect in mean field dynamo theory becomes proportional to a relaxation time scale multiplied by the difference between kinetic and current helicities. It is shown in \cite{BSM05} that the value of the relaxation time is positive and, in units of the turnover time at the forcing wavenumber, it is of the order of unity. Kinetic and current helicities are shown in \cite{BSM05} to be dominated by large scale properties of the flow. Recent studies in \cite{SSB07} of the nonlinear alpha effect showed that in the limit of small magnetic and hydrodynamic Reynolds numbers, both the second order correlation approximation (or first-order smoothing approximation) and the $\tau$ approximation give identical results. This is also supported by simulations \cite{BS07} of isotropically forced helical turbulence whereby the contributions to kinetic and magnetic alpha effects are computed. The study performed in \cite{BS07} provides an extra piece of evidence that the $\tau$ approximation is a useable formalism for describing simulation data and for predicting the behavior in situations that are not yet accessible to direct numerical simulations. | 7 | 10 | 0710.5052 |
0710 | 0710.0663_arXiv.txt | The intrinsic anisotropy $\delta$ and flattening $\epsilon$ of simulated merger remnants is compared with elliptical galaxies that have been observed by the SAURON collaboration, and that were analysed using axisymmetric Schwarzschild models. Collisionless binary mergers of stellar disks and disk mergers with an additional isothermal gas component, neglecting star formation cannot reproduce the observed trend $\delta = 0.55 \epsilon$ (SAURON relationship). An excellent fit of the SAURON relationship for flattened ellipticals with $\epsilon \geq 0.25$ is however found for merger simulations of disks with gas fractions $\geq 20\% $, including star formation and stellar energy feedback. Massive black hole feedback does not strongly affect this result. Subsequent dry merging of merger remnants however does not generate the slowly-rotating SAURON ellipticals which are characterized by low ellipticities $\epsilon < 0.25$ and low anisotropies. This indicates that at least some ellipticals on the red galaxy sequence did not form by binary mergers of disks or early-type galaxies. We show that stellar spheroids resulting from multiple, hierarchical mergers of star-bursting subunits in a cosmological context are in excellent agreement with the low ellipticities and anisotropies of the slowly rotating SAURON ellipticals and their observed trend of $\delta$ with $\epsilon$. The numerical simulations indicate that the SAURON relation might be a result of strong violent relaxation and phase mixing of multiple, kinematically cold stellar subunits with the angular momentum of the system determining its location on the relation. | A popular formation scenario for early-type galaxies is the collision and merger of two roughly equal-mass galaxies with mass ratios between 1:1 and 4:1. This famous major merger scenario \citep{1972ApJ...178..623T} has been very successful in explaining observed properties of ellipticals, like their kinematics, surface density profile or isophotal shape \citep{1981MNRAS.197..179G,1983MNRAS.205.1009N,1988ApJ...331..699B, 1990Natur.344..379B,1992ARA&A..30..705B,1992ApJ...400..460H, 1993ApJ...409..548H,1993A&A...278...23B,1999ApJ...523L.133N,2001ApJ...554..291C, 2003ApJ...597..893N,2005MNRAS.357..753G,2005MNRAS.360.1185J, 2005A&A...437...69B,2006MNRAS.369..625N,2006MNRAS.372..839N, 2006ApJ...641...21R,2006astro.ph..7446C}. Numerical simulations for example showed that the family of disky, fast rotating ellipticals could result from stellar disk galaxy mergers with unequal mass ratios of 3:1 to 4:1 \citep{1998giis.conf..275B,1998ApJ...502L.133B,1999ApJ...523L.133N, 2003ApJ...597..893N,2005A&A...437...69B,2006MNRAS.372..839N} or from gas-rich 1:1 to 2:1 disk mergers where the gas subsequently settles into the equatorial plane of the merger remnant and produces a new stellar disk component \citep{2005MNRAS.359.1379K, 1996ApJ...471..115B,2001ApJ...555L..91N,2002MNRAS.333..481B, 2005ApJ...622L...9S}. Boxy, slowly rotating ellipticals, on the other hand form in stellar disk-disk mergers with mass ratios of 1:1 to 2:1 \citep{1994ApJ...427..165H,2003ApJ...597..893N} or from multiple major disk mergers \citep{1996ApJ...460..101W}. \citet{2007arXiv0709.3439B} showed that repeated minor mergers result in remnant properties very similar to one corresponding to major mergers. A serious problem of the major merger scenario was the fact that cosmological models do not predict a dependence of the mass ratio of mergers on total galaxy mass or luminosity \citep{2003ApJ...597L.117K}. If, on the other hand, the mass ratio determines the isophotal shape and rotational properties of merger remnants one would expect that the ratio of the number of fast rotating, disky ellipticals to the number of slowly rotating, boxy systems should be independent of luminosity \citep{2006ApJ...636L..81N}. This is in contrast with observations that show a strong dependence of isophotal shape and rotational properties on galaxy mass. While massive galaxies are preferentially boxy, slow rotators, lower-mass ellipticals are predominantly disky and fast rotators (for a summary see \citealp{1996ApJ...464L.119K}). \citet{2003ApJ...597L.117K} argued that this mass dependence of galaxy properties could be explained as a result of differences in the morphologies of the merging progenitors (see e.g. \citealp{2007MNRAS.tmp..808K}). Using semi-analytical models they showed that gas-rich disk-disk mergers dominate at the low-mass end of ellipticals while intermediate mass ellipticals should have formed preferentially from mixed mergers involving a disk and an elliptical galaxy. Finally, the most massive early-type galaxies should have experienced a last elliptical-elliptical merger (dry merger) \citep{2006ApJ...636L..81N}. Mixed mergers have not yet been studied in details, although, according to \citet{2003ApJ...597L.117K} they should be more frequent than dry mergers (see however \citealp{2007arXiv0706.1243H}). Dry mergers and their implications for the formation of the red galaxy sequence have however received a lot of attention recently \citep{2005MNRAS.359.1379K,2005MNRAS.361.1043G,2005astro.ph..6044F, 2006ApJ...636L..81N,2006MNRAS.369.1081B}. Further refinement of theoretical models has recently been achieved by including black-hole physics in simulations of galaxy mergers and elliptical galaxy evolution. Energetic feedback from central black holes might solve some pending problems of major merger models, like the suppression of late inflow of cold gas and star formation that would make ellipticals look much bluer than observed \citep{2005ApJ...620L..79S}. In summary, despite several still unsolved questions (e.g. \citealp{2007astro.ph..2535N}), the major merger scenario has become a popular model in order to explain the origin of bulge-dominated, spheroidal galaxies (see e.g. \citealp{2007arXiv0706.1246H,2007arXiv0706.1243H}). Progress in understanding galactic evolution is often driven by strong interactions between observers and theorists. Increasingly more sophisticated theoretical/numerical models are confronted with continuously improving observations that lead to new theoretical challenges. One example is the SAURON project \citep{2004MNRAS.352..721E} which aims to determine the 2-dimensional structural and kinematical properties of early-type galaxies using a panoramic integral-field spectrograph. In order to interpret the observations and study the intrinsic galaxy structure, axisymmetric Schwarzschild models are applied to the observations \citep{2006MNRAS.366.1126C,2007MNRAS.379..418C}. The results, published so far, have revealed interesting fine structures and physical properties which provide new and deeper insight into the origin of galaxies, in particular when compared to simulation \citep{2000MNRAS.316..315B,2007MNRAS.376..997J}. In this paper we confront a recently published SAURON analysis of preferentially axisymmetric elliptical galaxies \citep{2007MNRAS.379..418C} with the predictions of numerical merger simulations and cosmological models of galaxy formation. Section 2 summarizes the observations. Section 3 shows that simulations of collisionless and gaseous disk-disk mergers, neglecting star formation and stellar feedback cannot reproduce the observational results. We demonstrate that star formation and stellar energy feedback has a strong effect on the final structure of merger remnants, leading to a good agreement with the SAURON observations of fast rotating ellipticals. The origin of the round, almost isotropic, slowly rotating SAURON ellipticals is explored in section 4. Isolated, dry mergers of ellipticals that formed as discussed in section 3 cannot explain these objects. We show however that cosmological initial conditions, leading to a series of multiple major and minor mergers, coupled with local star bursts generate spheroidal stellar systems in very good agreement with the observations. Conclusions follow in section 5. | The numerical simulations, discussed in the previous sections, have shown that interstellar gas dynamics, star formation and stellar feedback plays a crucial role in order to reproduce the observed kinematical and isophotal properties of fast rotating, early-type galaxies. The final structure of the merger remnants depends on the initial mass ratio and gas fraction. The remnants are more round, less anisotropic and more rotationally supported the smaller the mass ratio $M_1/M_2 \geq 1$ of the progenitors and the larger the initial gas fraction. The dependence of $\delta$ on $\epsilon_{int}$ is in agreement with the observed trend found in the SAURON sample. Subsequent dry re-merging of disk-disk merger remnants however does not generate the observed slowly-rotating red SAURON ellipticals with small anisotropies and ellipticities. This indicates that at least some early-type galaxies on the red galaxy sequence formed in a different way. We find that multiple mergers of stellar substructures that formed from cold gas infall into dark matter halos in cosmological simulations produce round, isotropic and slowly-rotating relaxed stellar systems that are in perfect agreement with the SAURON observations. Multiple mergers of stellar systems in dense group enviroments therefore appear to be a promising alternative scenario for the origin of the red, massive galaxy population. Despite the fact that merger simulations with star formation lead to a correlation between anisotropy and ellipticity ($\delta = 0.67 \times \epsilon_{int}$) that is very similar to that inferred from observations its origin is not understood yet. It is interesting that merger remnants appear to move closer towards this relation along lines of constant $V/\sigma$ (i.e. roughly constant specific angular momentum) in the case of a strong relaxation process. Here, strong relaxation is defined as the merger of a system of kinematically cold systems of stars that lateron break up and generate a kinematically hot stellar remnant. Several conditions could lead to such a violent dynamical process. The cold stellar clumps could for example have formed in the star-bursting tidal tails of interacting, gas-rich disk galaxies. Another possibility is the cosmological multiple merging of dark matter substructures with an embedded stellar systems. The SAURON relation might represent the relaxed and phase-mixed end state of these complex mergers with the location of the remnant on the relation being determined by its specific angular momentum which is related to its value of $V/\sigma$. More theoretical work will be required in order to better understand these interesting questions and their connection to early-type galaxy formation. | 7 | 10 | 0710.0663 |
0710 | 0710.2450_arXiv.txt | Smooth double crossing of the phantom barrier $w_{\Lambda} = -1$ has been found possible in cosmological model with Gauss-Bonnet-scalar interaction, in the presence of background cold dark matter. Such crossing has been observed to be a sufficiently late time phenomena and independent of the sign of Gauss-Bonnet-scalar interaction. The luminosity distance versus redshift curve shows a perfect fit with the $\Lambda CDM$ model up to $z=3.5$. | The puzzle associated with recent cosmic acceleration, triggered by $70\%$ of dark energy or more \cite{a1} is far from being resolved uniquely. In the mean time, cosmologists are being confronted with yet another more intriguing challenge to explain the crossing of the so called phantom divide line $(w_{\Lambda}=-1)$, at sufficiently late time of cosmological evolution. Some recent analysis \cite{a1},\cite{b1} of the presently available observational data are in favour of the value $w_{de}<-1$, at present., $w_{de}$ being the dark energy equation of state. There are also a lot of evidence all around \cite{c}, of a dynamical dark energy equation of state, which has crossed the so called phantom divide line $w_{\Lambda}=-1$ recently, at the value of red-shift parameter $z \approx 0.2$. Apparently though the problem turns out to be more serious and complicated, but then, the puzzle of crossing the phantom divide line has also rendered some sort of selection rule. $\Lambda CDM$-model, which is known to suffer from the disease of fine tuning (see \cite{d} for a comprehensive review) can now be ruled out due to the requirement of a dynamic state parameter. Further, if the analysis of Vikman \cite{e} is correct, then it is not possible to cross the phantom divide line in a single minimally coupled scalar field theory, without violating the stability both at the classical \cite{f} and also at the quantum mechanical levels \cite{g}, (though it has recently been inferred \cite{on} that quantum Effects which induce the $w<-1$ phase, are stable in the $\phi^4$ model). Thus single minimally coupled scalar field models like quintessence $(w>-1)$ and phantom $(w<-1)$ are to be kept aside. Consequently, we are now left with some what more complicated models. One of these is a hybrid model, composed of two scalar fields, viz, quintessence and phantom - usually dubbed as quintom model \cite{h}. Other models like non-minimal scalar tensor theory of gravity \cite{i}, hessence \cite{j} and models including higher order curvature invariant terms \cite{k} also exist in the literature. \\ Gauss-Bonnet term is yet another candidate which may be pursued for the purpose. The possibility of crossing the phantom divide line through Gauss-Bonnet interaction has been explored in some recent works \cite{l},\cite{m}. But then, these models are even complicated in the sense that either brane-world scenario \cite{l} or scalar field and matter coupling \cite{m} are invoked. In this article the possibility of smooth crossing of the phantom divide line $w_{\Lambda}=-1$ has been expatiated simply by introducing Gauss-Bonnet-Scalar coupling term in the 4-dimensional Einstein-Hilbert action.\\ Gauss-Bonnet term arises naturally as the leading order of the $\alpha'$ expansion of heterotic superstring theory, where, $\alpha'$ is the inverse string tension \cite{n}. Gauss-Bonnet term is topologically invariant and thus does not contribute to the field equations in four dimensions. However, the low energy limit of the string theory gives rise to the dilatonic scalar field which is found to be coupled with various curvature invariant terms \cite{o}. The leading quadratic correction gives rise to Gauss-Bonnet term with a dilatonic coupling \cite{p}. Therefore it is reasonable to consider Gauss-Bonnet interaction in four dimension with dilatonic-scalar coupling. Several works with Gauss-Bonnet-dilatonic coupling are already present in the literature \cite{q}. In particular, important issues like - late time dominance of dark energy after a scaling matter era and thus alleviating the coincidence problem, crossing the phantom divide line and compatibility with the observed spectrum of cosmic background radiation have also been addressed recently \cite{km}.\\ In a recent work with Gauss-Bonnet interaction \cite{a}, a solution in the form $a=a_{0}e^{A\sqrt t}$ ($a$ being the scale factor, and $A>0$) has is been found to satisfy the field equations with different forms (sum of exponentials, sum of inverse exponentials, sum of powers and even quadratic) of potentials. Solution in a more general form $(a=a_{0}e^{A t^f}), A>0, 0<f<1$, for Einstein's gravity with a minimally coupled scalar field was found in the nineties \cite{r} and was dubbed as intermediate inflation. We \cite{a}, on the other hand, observed that such solution depicts a transition from decelerated to accelerated expansion at sufficiently later epoch of cosmic evolution, which asymptotically goes over to de-Sitter expansion. Thus, it appeared that such solution may construct viable cosmological models of present interests. Under this consequence, a comprehensive analysis has been carried out \cite{b} with such solution in the context of a generalized k-essence model. It has been observed that it admits scaling solution with a natural exit from it at a later epoch of cosmic evolution, leading to late time acceleration with asymptotic de-Sitter expansion. The corresponding scalar field has also been found to behave as a tracker field \cite{s}, thus avoiding cosmic coincidence problem. \\ In the present work, we show that Gauss-Bonnet-Dilatonic scalar coupling with Einstein's gravity in four dimensions, admits solution in a general form $(a=a_{0}e^{A t^f}), A>0, 0<f<1$, which is viable of crossing the phantom divide line twice, once from above and the other from below in the recent epoch. Since the crossing is transient, so we may conclude that it does not show any pathological behaviour like Big-Rip \cite{f}, at least in the classical level. | Altogether we have obtained a late time transient crossing of the phantom divide line first from above and more recently from below, starting from the inclusion of a Gauss-Bonnet-dilatonic scalar coupling term in the standard Einstein-Hilbert action in four dimension. Since the crossing is transient, so such double crossing is free from any sort of pathological behaviour both at the classical \cite{f} and at the quantum mechanical \cite{g} levels. The striking feature of the model lyes in it's indistinguishability with the standard $\Lambda CDM$ model, in terms of the luminosity-redshift and more precisely for distance modulus-redshift curves. To identify between the two models we therefore require to observe dark energy equation of state $w_{de}$ independently. If $w_{de}$ is truly found to be dynamical and has really encountered a recent crossing of the phantom divide line, then only we can definitely distinguish the standard $\Lambda CDM$ model with the present one. It is highly interesting to learn that smooth transient crossing of the phantom divide line is allowed for both negative and positive $(\lambda \lessgtr 0)$ type of Gauss-Bonnet-scalar interaction. Figures (7), (8) and (9) also reveal that it is true even for different forms of the potential. Such transient crossing independent of the value of $\lambda$ also signals that it might be possible to carry out the same treatment even for a single scalar field model. \textbf{Acknowledgement}:{Acknowledgement is due to Dipartimento di Scienze Fisiche, Universit$\grave{a}$ degli studi di Napoli, Federico II, for their hospitality, TRIL (ICTP) for financial assistance and Prof. Claudio Rubano for some illuminating discussion.} | 7 | 10 | 0710.2450 |
0710 | 0710.1948_arXiv.txt | {The primary proton spectrum up to 100 TeV has been investigated by balloon- and satellite-borne instruments. Above this energy range only ground-based air shower arrays can measure the cosmic ray spectrum with a technique moderately sensitive to nuclear composition. An array which exploits the full coverage approach at very high altitude can achieve an energy threshold well below the TeV region, thus allowing, in principle, the inter-calibration of the measured proton content in the primary cosmic ray flux with the existing direct measurements from balloons/satellites. The capability of the ARGO-YBJ experiment, located at the YangBaJing Cosmic Ray Laboratory (4300 m a.s.l., Tibet, P.R. China), in selecting the surviving primary cosmic ray protons is discussed. A procedure looking for quasi-unaccompanied events with a very steep lateral distribution is also presented.} \begin{document} | Despite large progresses in operating new multi-component Extensive Air Shower (EAS) experiments and in the analysis techniques to infer energy spectra and chemical composition, the key questions concerning the origin of the "knee" in the cosmic ray energy spectrum are still open. In particular, one of the most important questions to be solved is the position of the proton knee (see for example \cite{pds} and reference therein). In fact, different experiments have claimed to see a proton knee at different energies: at 10 TeV the MUBEE collaboration \cite{mubee}, at a few hundreds TeV the TIBET-AS$\gamma$ experiment \cite{tibet} and at a few PeV the KASCADE \cite{kascade} and EAS-TOP \cite{eastop} experiments. In addition, direct measurements carried out in the 100 TeV region by RUNJOB \cite{runjob} and JACEE \cite{jacee} do not exhibit any spectral break up to the highest measured energy ($\sim$ 800 TeV). The knowledge of the primary proton spectrum is of great importance to understand the cosmic rays acceleration mechanisms and propagation processes in the Galaxy. A careful measurement of the proton spectrum over a wide range of primary energies (from 0.1 TeV to 10 PeV) using the same method is one of the main tasks of the future cosmic ray experiments. The energy region up to about 100 TeV has been investigated by balloon- and satellite-borne instruments, and above these energies only ground-based air shower experiments may provide data which have however poor mass resolution. The detection of single hadrons at ground level, strongly related to the surviving primary cosmic ray protons, has been recognized for a long time as a method to investigate the proton spectrum over a large energy range and to derive the total inelastic cross section \cite{yodh,eastop-psgl,kascade-psgl}. The ARGO-YBJ experiment, an air shower array exploiting the full coverage approach at the YangBaJing Cosmic Ray Laboratory (Tibet, P.R. China, 4300 m a.s.l., 606 g/cm$^2$), offers the unique opportunity to investigate the cosmic ray spectrum over a large energy range (about 3 decades) because of its ability to operate down to a few TeV, thus overlapping the direct measurements, by measuring small size air showers (strip or digital read-out with high spatial granularity) and up to the PeV region by measuring the RPCs charge (analog read-out \cite{argo_bigpad}). In this paper we will discuss how a direct measurement of the primary proton spectrum could be achieved up to the PeV energies with the ARGO-YBJ experiment. | The selection of surviving primary protons is possible with the ARGO-YBJ experiment by measuring events interacting a few g/cm$^2$ above the observational level. The granularity of the detector allows to investigate in great details the lateral and temporal features of hadronic jets just produced. Calculations are in progress to evaluate the contribution of heavier primaries and to properly take into account the fluctuations in the shower development. A preliminary data analysis will be presented at the conference. | 7 | 10 | 0710.1948 |
0710 | 0710.3984_arXiv.txt | We argue that giant flares in SGRs can be associated to the core conversion of an isolated neutron star having a subcritical magnetic field $\sim 10^{12}$ G and a fallback disk around it. We show that, in a timescale of $\lesssim 10^5$ yrs, accretion from the fallback disk can increase the mass of the central object up to the critical mass for the conversion of the core of the star into quark matter. A small fraction of the neutrino-antineutrino emission from the just-converted quark-matter hot core annihilates into $e^+e^-$ pairs above the neutron star surface originating the gamma emission of the spike while the further cooling of the heated neutron star envelope originates the tail of the burst. We show that several characteristics of the giant flare of the SGR 1806-20 of 27 December 2004 (spike and tail energies, timescales, and spectra) can be explained by this mechanism. | 7 | 10 | 0710.3984 |
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0710 | 0710.4871_arXiv.txt | We present the detection of low-amplitude, long-period $g$-modes in two individual sdBV stars which are known to be $p$-mode pulsators. Only few of these hybrid objects, showing both $p$- and $g$-modes, are known today. We resolve the $g$-mode domain in HS\,0702+6043 and add HS\,2201+2610 to the list of hybrid pulsators. To discover the low-amplitude $g$-modes, a filtering algorithm based on wavelet transformations was applied to denoise observational data. | Subdwarf B stars are evolved and compact objects which are thought to be core helium burning. They populate the extreme horizontal branch (EHB) at effective temperatures of $20\,000$\,--\,$40\,000$~K and surface gravities $\log (g/\rm{cm\,s^{-2}})$ between 5.0 and 6.2. Variable sdB stars (sdBVs) can be divided into the following subclasses: The $p$-mode pulsators show higher amplitudes and shorter periods at higher temperatures. \citet{rl_kil97} discovered EC 14026-2647 as the class prototype. The $g$-mode pulsators have lower amplitudes and longer periods at lower temperatures and the class prototype PG\,1716+426 was discovered by \citet{rl_gre03}. Hybrid pulsators show both $p$- and $g$-modes. These are particularly exciting objects, since the two mode types probe different regions in the star. Both mode types are thought to be driven by the same $\kappa$-mechanism \citep{rl_cha97,rl_fon03,rl_jes06}. Figure 1 shows the location of some sdB pulsators in the $\log g$\,--\,$T_{\rm eff}$ diagram. \begin{figure}[!ht] \centering \includegraphics[scale=0.45,angle=-270]{loggteff_proceedings_color.eps} \caption{Location of some known sdB pulsators in the $\log g$-$T_{\rm eff}$ diagram. Evolutionary tracks \citep{rl_dor93} for three stellar masses are indicated, as well as helium main sequence (HeMS) and zero age (ZAEHB) and terminal age (TAEHB) extreme horizontal branches. The positions (taken from E. M. Green and G. Fontaine; for a discussion see also Green et al., these proceedings) clearly reveal two instability regions. The three known hybrid pulsators are labelled 1, 2 and 3.} \end{figure} | 7 | 10 | 0710.4871 |
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0710 | 0710.3383_arXiv.txt | As part of our on-going investigation of magnetic activity in ultracool dwarfs we present simultaneous radio, X-ray, UV, and optical observations of \lsr\ (M8.5), and simultaneous X-ray and UV observations of \vb\ (M8), both with a duration of about 9 hours. \lsr\ exhibits persistent radio emission and H$\alpha$ variability on timescales of $\sim 0.5-2$ hr. The detected UV flux is consistent with photospheric emission, and no X-ray emission is detected to a deep limit of $L_X/L_{\rm bol}\lesssim 10^{-5.7}$. The H$\alpha$ and radio emission are temporally uncorrelated, and the ratio of radio to X-ray luminosity exceeds the correlation seen in F--M6 stars by $\gtrsim\!2\times 10^4$. Similarly, $L_{\rm H\alpha}/L_X\gtrsim 10$ is at least 30 times larger than in early M dwarfs, and eliminates coronal emission as the source of chromospheric heating. The lack of radio variability during four rotations of \lsr\ requires a uniform stellar-scale field of $\sim 10$ G, and indicates that the H$\alpha$ variability is dominated by much smaller scales, $\lesssim 10\%$ of the chromospheric volume. \vb, on the other hand, shows correlated flaring and quiescent X-ray and UV emission, similar to the behavior of early M dwarfs. Delayed and densely-sampled optical spectra exhibit a similar range of variability amplitudes and timescales to those seen in the X-rays and UV, with $L_{\rm H\alpha}/L_X\sim 1$. Along with our previous observations of the M8.5 dwarf \tvlm\ we conclude that late M dwarfs exhibit a mix of activity patterns, which points to a transition in the structure and heating of the outer atmosphere by large-scale magnetic fields. We find that rotation may play a role in generating the fields as evidenced by a tentative correlation between radio activity and rotation velocity. The X-ray emission, however, shows evidence for super-saturation at $v{\rm sin}i\gtrsim 25$ km s$^{-1}$, which could be the result of secondary effects such as inefficient heating or centrifugal stripping of extended coronal loops. These effects may underlie the severe violation of the radio/X-ray correlation in ultracool dwarfs. Upcoming observations of L dwarfs will reveal whether these trends continue in substellar objects. | \label{sec:intro} Over the past several decades, observations of magnetic activity in stars of spectral type F--M have uncovered a variety of correlations between the different activity indicators, as well as between the level of activity and properties such as stellar rotation and age. Coronal X-ray emission, transition region UV emission, chromospheric H$\alpha$ emission, and non-thermal radio emission increase with both rotation and youth (e.g., \citealt{kra67,pgr+81,sis+88}), and are temporally and energetically correlated in quiescence and during flares (e.g., \citealt{cra82,gb93,bg94,hgr96}). These observations have led to a general paradigm of magnetic field amplification at the shearing interface between the radiative and convective zones -- the so-called $\alpha\Omega$ dynamo \citep{par55}. This dynamo operates through a combination of stretching by differential rotation ($\Omega$) and twisting by convective motions ($\alpha$). In quiescence, the dynamo-generated fields provide a heating source for the chromospheres and coronae primarily through magnetic waves and field dissipation on small scales (e.g., \citealt{nu96,apr01}). The interplay between the chromosphere, transition region, and corona is not fully understood, but is known to involve a complex combination of radiation, conduction, and mass flows. In dMe stars it has been proposed that the transition region and chromosphere are instead heated by coronal X-rays, leading to the observed typical luminosity ratios of $L_{\rm CIV}/L_X\sim 10^{-1.5}$ and $L_{\rm H\alpha}/L_X\sim 10^{-0.5}$, respectively (e.g., \citealt{cra82,hj03}). In addition to quiescent emission, sudden and large-scale dissipation of the field through large-scale magnetic reconnection may lead to particle acceleration and evaporation of the lower atmosphere into the chromosphere, transition region, and corona, giving rise to correlated radio, X-ray, UV, and H$\alpha$ flares (e.g., \citealt{neu68}). The level of activity increases with faster rotation, but eventually saturates at $L_X/L_{\rm bol}\sim 10^{-3}$ and $L_{\rm H\alpha}/L_{\rm bol}\approx 10^{-3.5}$ for rotation periods of $\lesssim 3$ d in F--M stars. It remains unclear whether this trend reflects a saturation of the dynamo itself or secondary centrifugal effects such as coronal stripping or sweeping of the field toward the poles (e.g., \citealt{vil84,ju99,ssv01}). These activity trends continue to hold even beyond spectral type $\sim {\rm M3}$, where the stellar interiors become fully convective and the $\alpha\Omega$ dynamo can no longer operate. Indeed, the level of X-ray and H$\alpha$ activity peaks in mid M dwarfs at saturated levels \citep{vw87,gmr+00,mb03,whw+04}. This suggests that whatever dynamo mechanism operates in these low mass stars, it is already present in higher mass objects, and becomes increasingly dominant beyond spectral type M3. However, the level of X-ray and H$\alpha$ activity drops precipitously beyond spectral type M7, reaching mostly non-detectable levels by spectral type L5. This decrease is accompanied by a clear transition from quiescent emission to flares in the few percent of active objects \citep{rkg+99,gmr+00,rbm+00,lkc+03,whw+04}. These trends point to a change in the dynamo mechanism, the field configuration, and/or the field dissipation process. Radio observations, however, have uncovered a substantial fraction of active late M and L dwarfs ($\sim 10\%$ or higher; \citealt{ber06}), which exhibit both quiescent and flaring emission \citep{bbb+01,ber02,brr+05,bp05,ber06,ohb+06,adh+07,pol+07,hbl+07,aob+07,bgg+07}. Unlike the trend in H$\alpha$ and X-rays, the level of radio activity appears to {\it increase} with later spectral type \citep{ber02,ber06}. The radio emission is orders of magnitude brighter than expected based on the radio/X-ray correlations that are observed in stars down to spectral type M6, and requires magnetic fields of $\sim 0.1-3$ kG with order unity filling factor sustained over timescales of at least several years \citep{brr+05,ber06,bgg+07}. Thus, contrary to evidence from X-rays and H$\alpha$, a substantial fraction of ultracool dwarfs continue to generate and dissipate magnetic fields. In order to investigate the field generation and dissipation in detail, we have recently initiated a program of simultaneous, multi-wavelength observations of ultracool dwarfs. Such observations are required to trace the temporal evolution of flares across the corona, transition region, and chromosphere, and to study the relation between particle acceleration and heating, particularly in the context of the known correlations in F--M stars. In a previous paper (\citealt{bgg+07}; hereafter Paper I) we presented observations of the M8.5 rapid rotator \tvlm, which exhibited a wide range of temporally uncorrelated emission in the radio, H$\alpha$, and X-rays. These included quiescent radio emission from a large-scale field, radio flares from a tangled field component with $B\approx 3$ kG, and periodic H$\alpha$ emission matching the stellar rotation period with an inferred hot spot covering about $50\%$ of the stellar photosphere. In addition, the quiescent H$\alpha$ emission exceeded the X-ray luminosity by about a factor of two, likely ruling out coronal X-ray emission as the source of chromospheric heating. Here we present observations of the M8 and M8.5 dwarfs \vb\ and \lsr, both of which are known to exhibit magnetic activity. We find substantially different behavior in each of these two objects, with correlated X-ray/UV flares and quiescent emission in \vb, and uncorrelated radio/H$\alpha$ emission in \lsr. Along with \tvlm, the mixed behavior in late M dwarfs thus indicates a transition in the properties of the magnetic field and its dissipation in this spectral type range. We show that rotation may play at least a partial role in explaining these trends. | \label{sec:conc} We presented simultaneous multi-wavelength observations of two late M dwarfs to trace the magnetic activity in their outer atmospheres. In the case of \lsr\ (M8.5) we detect persistent radio emission and highly variable H$\alpha$ emission. No excess UV emission is detected. Similarly, we detect no X-ray emission to a deep limit of $L_X/L_{\rm bol}\lesssim 10^{-5.7}$. The ratio of radio to X-ray flux exceeds the average value measured in a wide range of active stars by more than 4 orders of magnitude. Similarly, the ratio $L_{\rm H\alpha}/L_X\gtrsim 10$ exceeds the values measured for M0--M6 dwarfs by at least an order of magnitude, and rules out heating of the quiescent chromosphere by coronal emission. Temporally, we find no correspondence between the radio and H$\alpha$ light curves. This indicates that even if the source of chromospheric heating is magnetic reconnection, it occurs on sufficiently small scales that the overall radio emission does not change by more than $\sim 10\%$. Alternatively, the radio emission produced in conjunction with the H$\alpha$ variability may peak at a lower frequency than our 8.5 GHz observations. In the former scenario, our limit on the radio variability during the brightest H$\alpha$ episode can be interpreted as a limit of $\lesssim 10\%$ on the chromospheric volume involved in the variable H$\alpha$ emission. In the latter scenario, it is possible that the coincident radio emission is dominated by coherent emission at $\nu\approx 2.8\times 10^6\,B$, with $B\lesssim 3$ kG. With our measured rotation velocity of $v{\rm sin}i\approx 50$ km s$^{-1}$ for \lsr, the radio observations sample nearly 4 full rotations. The stability of the radio emission thus requires a large-scale and uniform magnetic field, with an inferred strength of $\sim 10$ G. The similarity in flux to a past observation suggests that the field is stable on year timescales. This may be expected given the long convective turnover time for ultracool dwarfs, on the order of several years. \vb, on the other hand, exhibits bright and variable X-ray and UV emission, including a pair of flares and quiescent emission. The behavior in both bands, which trace the corona and transition region, respectively, is highly correlated. From the X-ray spectrum we infer a coronal temperature of about $3\times 10^6$ K, with a possible second component at $\sim 10^7$ K. The optical emission lines exhibit a similar amplitude and timescale of variability to the X-rays and UV, with two distinct states of gradual and impulsive flaring. The lack of causal relation between the two states points to emission from distinct plasma environments. Moreover, the shallow Balmer decrement and \ion{He}{1} emission in the impulsive Flare state require significantly denser and hotter plasma. Our detailed study of \lsr, \vb, and \tvlm\ reveals mixed patterns of of behavior compared to the magnetic activity in F-M6 stars, and thus points to a transition in the atmospheric structure and heating process in the late M dwarf regime. In particular, stellar-scale magnetic fields are present, as evidence by quiescent and uniform radio emission, but the coronae are generally weaker than in early M dwarfs. The chromospheric activity declines less rapidly, and all three targets exhibit a similar range of highly variable H$\alpha$ emission, $L_{\rm H\alpha}/L_{\rm bol}\sim 10^{-5}-10^{-4.5}$. This range is about an order of magnitude less that the saturated value of mid M dwarfs. In the standard picture of solar and stellar magnetic flares, the release of magnetic energy through processes such as reconnection leads to acceleration of electrons, and subsequently heating of the chromosphere, transition region, and corona through evaporation of the lower atmosphere. In quiescence, the chromospheres of M0--M6 dwarfs may instead be heated by coronal X-ray emission \citep{cra82}, as evidenced by the typical observed ratios of $L_{\rm H\alpha}/L_X\sim 1/3$ \citep{cra82,hgr96}. Such a mechanism is clearly not at play in the case of \lsr, and most likely \tvlm, since the chromospheric emission is significantly more energetic than the non-detected corona. It is possible instead that the source of chromospheric heating is continuous micro-flaring activity, leading to replenishment of chromospheric and transition region plasma by evaporation. In this scenario the weak or absent coronal emission may be due to a limited temperature enhancement of $\lesssim 10^5$ K, which is sufficient to produce chromospheric emission, but no significant soft X-ray emission. The continuously variable H$\alpha$ emission in \lsr\ appears to support this idea, and along with the lack of corresponding radio variability points to heating on scales much smaller than the stellar photosphere. As we noted in \S\ref{sec:rot}, it is also possible that rapid rotation in ultracool dwarfs suppresses the X-ray emission through centrifugal stripping. The apparent increase in radio activity with rotation indicates that the dynamo itself does not saturate, at least up to $v{\rm sin}i\sim 60$ km s$^{-1}$, which is roughly $1/3$ of the break-up velocity. The effects leading to the apparent super-saturation in the X-rays may also underlie the severe violation of the radio/X-ray correlation since the level of violation appears to be correlated with rotation velocity (Figure~\ref{fig:bgrot}). We stress that the role of rotation is still speculative due to the small number of objects with X-ray, radio, and rotation measurements. It is therefore crucial to increase the sample of ultracool dwarfs for which these quantities are measured. To summarize, we conclude that late M dwarfs mark a change in the properties of the magnetic field and its dissipation, as well as the generation of high temperature plasma in the outer atmosphere. The weak X-ray emission, both in relation to $L_{\rm bol}$ and $L_{\rm H\alpha}$, points to inefficient heating of plasma to coronal temperatures, or alternatively to stripping of the most extended magnetic loops. The lack of temporal correlation between the radio and H$\alpha$ activity, however, indicates that even if atmospheric evaporation is taking place, it occurs on smaller physical scale than the overall magnetic field structure. We end by noting that the use of simultaneous, multi-wavelength, and long duration observations to probe the magnetic activity of individual ultracool dwarfs provides a crucial complement to large single-band surveys. In particular, the long duration and high time resolution data elucidate the range of timescales and amplitudes of gradual and impulsive flares, and can moreover uncover rotationally-induced modulations, as in the case of \tvlm\ (Paper I). In addition, the existence or absence of correlations between the various activity indicators (as in \lsr\ and \tvlm), and their implications for the magnetic field and its dissipation, could not have been inferred from existing observations. In the upcoming Chandra cycle we will expand our analysis with observations of several L0--L3 dwarfs. These observations will reveal whether the transition in magnetic activity patterns seen in the late M dwarfs continues to lower mass objects, some of which are bona-fide brown dwarfs. | 7 | 10 | 0710.3383 |
0710 | 0710.2336_arXiv.txt | { We derive an expression for the entropy of a present dark matter halo described by a Navarro-Frenk-White modified model with a central core. The comparison of this entropy with the one of the halo at the freeze-out era allows us to obtain an expression for the relic abundance of neutralinos, which in turn is used to constrain the parameter space in mSUGRA models, when used with the WMAP observations. Moreover, by joning these results with the ones obtained from the usual abundance criteria, we are able to clearly discriminate validity regions among $\tb$ values of the mSUGRA model, by demanding both criteria to be consistent with the 2 sigma bounds of the WMAP observations for the relic density: $0.112<\Omega h^2<0.122$. We found that for $sgn~ \mu =+$, small values of $\tb$ are not favored; only for $\tb \sim 50$ are both criteria significantly consistent. The use of both criteria also allows us to put a lower bound on the neutralino mass, $m_{\chi}\geq151$GeV. } \PACS{ {14.80.Ly}{Supersymmetric partners of known particles} \and {95.35.+d}{Dark matter} \and {98.62.Gq}{Galactic halos} } % | s^h-s^f=ln\left[\frac{n_{\chi}^f}{n_{\chi}^h}\left(\frac{x^f}{x^h}\right)^{3/2}\right]~, \end{equation} where $x=m_{\chi}/T$, $T$ is the temperature of the gas. A region that fits with the conditions associated with the intermediate scale is the central region of halos ($~10 pc^3$ within the halo core); evaluating the thermodynamical quantities at this region, using equation (\ref{entropy}) and some extra assumptions (conservation of photon entropy), it is possible to construct a theoretical estimate for $s^h$ that depends on the nature of neutralinos ($m_{\chi}$ and $\langle\sigma v\rangle$), initial conditions (given by $x^f$), cosmological parameters ($\Omega_{\chi}$, the Hubble parameter, $h$) and structural parameters of the virialized halo (central values for temperature and density); for details of these and the following, see section IV of \cite{Cabral-Rosetti:2004kd}. An alternative estimate for $s^h$ can be made based on empirical quantities for observed structures in the present Universe using the microcanonical entropy definition in terms of phase space volume, but restricting this volume to the actual range of velocities accessible to the central particles. That is, restricting the escape velocity up to a maximal value $v_e(0)$ which is related to the central velocity dispersion of the halo ($\sigma_h$) by an intrinsic parameter $\alpha$: $v_e^2(0)\sim\alpha\sigma_h^2(0)$. In a recent work \cite{nuevo}, we estimate the value of $\alpha$ using an NFW modified model with a central core, and obtain $16.4<\alpha <27.8$. The range of values allowed for this parameter is of the highest importance to determine the allowed region of the parameter space in the mSUGRA model as will be clear in the results presented on next section. Equating the theoretical an empirical estimates for the entropy per particle it is obtained a relation for the relic abundance of neutralinos using the EC criterion\footnote{This formula is a small modification to the one presented in \cite{Cabral-Rosetti:2004kd}}: \begin{equation}\label{consistency} ln(\Omega_{\chi}h^2)=10.853-x^f+ln\left[\frac{(x^f\alpha)^{3/2}m_{\chi}}{f_g^*(x^f)}\right] \end{equation} where $f_g^*(x^f)$ is a function related to the degrees of freedom at the ``freeze-out'' time (see for example \cite{gondolo}) that will be described elsewhere \cite{nuevo}. Modifying the program micrOMEGAs, we obtained the value for $x^f$ for any region of the parameter space and then $\Omega_{\chi}$ using Eq.~(\ref{consistency}), therefore we were able to discriminate regions that are consistent with the WMAP constraints for the EC criterion. | Using both the AC and EC that have been described in the preceeding sections, we can compute the total mass density of neutralinos present today and constrain the region in the mSUGRA parameter space where both criteria are fulfilled. Out of the five parameters, we will fix $\mu >0$. Then our strategy is to explore wide regions for the values of the other four parameters, by means of a bi-dimensional analysis in the $m_0 - m_{1/2}$ plane for different fixed values of $A_0$ and tan$\beta$. It is important to mention that we are not presenting an exhaustive search in all the possible regions, but we concentrated on those regions which have received more attention in the literature, see for example \cite{belanger}. In Fig.~(\ref{tb10new}), we present the results for tan$\beta=10$, for three values of $A_0$, namely $A_0=1000, 0, -1000$ GeV, shown in the top, middle and bottom panels respectiveley. The yellow region (lower right corner) is where the $\tilde{\tau}$ is the LSP, the lighter and darker areas (red and blue for the online version in colours) define the allowed regions for the EC and AC respectively according to the observed DM density. The area of the EC region depends on the size of the interval of values of the parameter $\alpha$, the lower and upper bounds of $\alpha$ determine the upper and lower boundaries of the EC region. As can be seen from the figure, the region where both criteria are fullfilled is very small, in fact, only for the highest values of $\alpha$ there is an intersection between both criteria. This behavior holds for all values of $A_0$ in the interval $[-1000, 1000]$ GeV, here we are showing only the extreme and central values. \begin{figure}\centering \includegraphics[height=15cm, width=4.8cm, viewport=200 0 300 350]{mytb_10.pdf} \caption{Allowed regions in the parameter space for AC (lighter gray/red) and EC (darker grey/blue) for the mSUGRA model with sgn$\mu=+$, tan$\beta=10$, and $A_0=1000$ GeV, top panel, $A_0=0$ GeV, middle panel, and $A_0=-1000$ GeV, bottom panel. The figures show the so called bulk and coannihilation regions. The yellow region shows where the stau is the LSP.} \label{tb10new} \end{figure} Repeating the same procedure for larger values of tan$\beta$, it is found that the intersection region for both criteria becomes larger, but it gets to be significant for the largest values of this parameter. This is clearly shown in Fig.~(\ref{tb50new}), which is equivalent to Fig.~(\ref{tb10new}), but for tan$\beta=50$. In this case the bottom panel is for $A_0=-500$ GeV. It is clear from the figure that for these values of tan$\beta$ both criteria are consistent, as shown by the large intersection area for values of $A_0$ in the interval $[0,1000]$ GeV. For negative values of $A_0$ the intersection region decreases with $A_0$, see the bottom panel of the figure. For even lower values of $A_0$ the intersection becomes insignificant. \begin{figure}\centering \includegraphics[height=15cm, width=4.8cm, viewport=200 0 300 350]{mytb_50.pdf} \caption{The same as Fig.~(\ref{tb10new}), but for tan$\beta=50$, and now $A_0=-500$ GeV in the bottom panel.} \label{tb50new} \end{figure} In Figs.~(\ref{results2}) we present the same analysis but for the Focus Point region, and for the central value $A_0=0$. The situation is consistent with the previous results, both criteria intersect for $\tan \beta=50$ and there is nearly no intersection for $\tan \beta=10$. \begin{figure}\centering \includegraphics[width=8cm]{fig3.pdf} \includegraphics[width=8cm]{fig4.pdf} \caption{Allowed regions in the parameter space for AC (lighter gray/red) and EC (darker grey/blue) in the mSUGRA model with $A_0=0$, sgn$\mu=+$, tan$\beta=10$, top panel, and tan$\beta=50$, bottom panel. The region shows the so called Focus Point region.} \label{results2} \end{figure} This analysis allows us to arrive to one of the main results of our work. The use of both criteria favours large values of tan$\beta$. In Figs.(\ref{fig:Mhiggsvsm12}) and (\ref{fig:chi-higgs1}) we show the allowed values for the LSP and the Higgs mass after constraining the parameter space with the abundance and entropy criteria. As can be seen from Fig. (\ref{fig:Mhiggsvsm12}), the current limit for the Higgs favour, combined with the AC and EC criteria, favours even more a large value of $\tan \beta$. This, in turn, puts a constraint on the allowed SUSY mass spectra of the bulk and coannihilation regions: it gives an LSP of mass $m_{\chi} \sim 140$ GeV for $\tb ~ 10$, and a lower bound for the LSP mass $m_{\chi} \gtrsim 150$ GeV for large $\tb$. \begin{figure}\centering \includegraphics[width=8cm]{mhiggsvsm12.pdf} \caption{Allowed values for $M_{Higgs}$ as function of $m_{1/2}$. As can be seen from the figure, the present bound on the Higgs mass favours a large $\tan \beta$.} \label{fig:Mhiggsvsm12} \end{figure} \begin{figure}\centering \centerline{\includegraphics[width=8cm,angle=0]{fig6.pdf}} \caption{The figure shows the LSP mass plotted versus the Higgs mass, points above the dashed line are the allowed values for the LSP. The points in blue correspond to the Focus Point region, the ones in red to the bulk and coannihilation regions.} \label{fig:chi-higgs1} \end{figure} Further analysis, which is currently under way, is required to give more precise conclusions about this new method to constrain the parameter space of the mSUGRA model \cite{nuevo}. \vspace{0.3cm} We acknowledge partial support by CONACyT\\ M\'exico, under grants 32138-E, 34407-E and 42026-F, and PAPIIT-UNAM IN-122002, IN117803 and \\ IN115207 grants. JZ acknowledges support from DGEP-UNAM and CONACyT scholarships. | 7 | 10 | 0710.2336 |
0710 | 0710.0619_arXiv.txt | The large majority of EGRET point sources remain without an identified low-energy counterpart, and a large fraction of these sources are most likely extragalactic. Whatever the nature of the extragalactic EGRET unidentified sources, faint unresolved objects of the same class must have a contribution to the diffuse extragalactic gamma-ray background (EGRB). Understanding this component of the EGRB, along with other guaranteed contributions from known sources, is essential if we are to use this emission to constrain exotic high-energy physics. Here, we follow an empirical approach to estimate whether a potential contribution of unidentified sources to the EGRB is likely to be important, and we find that it is. Additionally, we show how upcoming GLAST observations of EGRET unidentified sources, as well as of their fainter counterparts, can be combined with GLAST observations of the Galactic and extragalactic diffuse backgrounds to shed light on the nature of the EGRET unidentified sources even without any positional association of such sources with low-energy counterparts. | \label{intro} The origin of the isotropic diffuse emission (Sreekumar et al.\ 1998) in energies between $100 {\, \rm MeV}$ and $20 {\, \rm GeV}$, detected by the {\it Energetic Gamma-Ray Experiment Telescope} (EGRET) aboard the {\it Compton Gamma-Ray Observatory}, remains one of the great unknowns of GeV-energy astrophysics. There are two major questions that still remain unanswered. 1.\ How much of the diffuse emission detected by EGRET is, in fact, extragalactic, and what is the spectrum of this extragalactic background? And 2.\ what fraction of the extragalactic emission can be attributed to each of the observationally established classes of gamma-ray emitters? Despite the associated uncertainties, these two issues are critical in any attempt to use gamma-ray observations to constrain exotic high-energy physics and yet-undetected classes of theorized gamma-ray emitters. To answer the first question, a good understanding of the Galactic diffuse emission in the EGRET energy range is required. In order to obtain the intensity and spectrum of the extragalactic emission from the EGRET sky maps, the Galactic emission needs to be modeled and subtracted. This is made complicated by the discrepancy between the observed Milky Way spectrum in energies of $\gtrsim 1 {\rm GeV}$ and theoretical expectations (Hunter et al.\ 1997). The observed spectrum is more shallow than model predictions based on the local demodulated cosmic ray spectrum. This deviation is known as the ``GeV excess'', and although various explanations have been proposed to account for part or all of the discrepancy, its origin remains a matter of debate (e.g., Pohl et al.\ 1997; B\"{u}sching et al.\ 2001; Strong et al.\ 2004b; Kamae et al.\ 2005; de Boer et al.\ 2006; Strong 2006; Stecker \etal 2007). As a result, determinations of the gamma-ray background using different Galactic emission models yield very different answers, both in intensity and in spectrum, despite being based on the same set of observations (e.g., Sreekumar \etal 1998; Strong et al.\ 2004a). Attempts to answer the second question have been plagued by uncertainties in the cosmic density and evolution of the two established classes of extragalactic gamma-ray emitters: normal galaxies and blazars. Our observational knowledge of the gamma-ray properties of normal galaxies is very limited, as the sample of normal galaxies which have been observed in gamma rays consists of only the Milky Way and a marginal detection of the Large Magellanic Cloud (Sreekumar et al.\ 1992; Hunter et al.\ 1997; Hartman et al.\ 1999). For this reason, the accuracy of theoretical estimates of the contribution of normal galaxies to the gamma-ray background is unavoidably at the order-of-magnitude level (e.g., Lichti et al.\ 1978; Pavlidou \& Fields 2002). But even in the case of blazars, which are by far the most numerous and best studied class of identified gamma-ray emitters, estimates of their contribution to the gamma-ray background vary from a few percent to $100\%$ of the background originally reported by the EGRET team (e.g., Padovani \etal 1993; Stecker \& Salamon 1996a; Kazanas \& Perlman 1997; Mukherjee \& Chiang 1999; M\"{ucke} \& Pohl 2000; Narumoto \& Totani 2006; Dermer 2007). The issue is further complicated by the existence of 171 sources which, at the time of publication of the 3rd EGRET catalog (hereafter 3EG; Hartman \etal 1999), had not been positively or potentially associated with a lower-energy counterpart. These sources are collectively known as the unidentified EGRET sources, and they are more numerous than any established group of gamma-ray emitters. The distribution of these sources on the sky is such that a Galactic feature can be clearly distinguished - however a large number of sources are located away from the Galactic plane and the Galactic center\footnote{Note however that the presence of sources at high latitudes does not, in itself, constitute proof that these objects are extragalactic (see e.g.\ the Gould Belt discussion in Gehrels et al.\ 2000).}. No more than a handful of sources can be associated with the Milky Way halo if the Milky Way is not many times brighter in gamma-rays than similar galaxies such as M31 (Siegal-Gaskins \etal 2007). Hence, it is almost certain that the EGRET unidentified sources include a significant extragalactic component. Although the nature of these extragalactic sources remains unknown, it is reasonable to believe that there is a large number of fainter, unresolved objects of the same class, which are guaranteed to have {\em some} contribution to the extragalactic gamma-ray background (EGRB). If these sources represent yet unidentified members of some known class of gamma-ray emitters (e.g.\ blazars), then excluding them from any calculation of the contribution of the parent class to the diffuse background would lead to a significantly underestimated result due to an incorrect normalization of the bright-end of the gamma-ray luminosity function. If they represent an unknown class of gamma-ray emitters, then the contribution of their unresolved counterparts to the diffuse emission would significantly limit the diffuse flux left to be attributed to known classes, exotic processes, and truly diffuse emission. Hence, some contribution of unresolved unidentified sources to the EGRB is certain. It is therefore clear that until we either answer the question of the nature of unidentified sources or derive some strong constraint indicating that a possible contribution of such unresolved objects to the EGRB would indeed be minor, we cannot hope to fully understand the origin of the EGRB. Detailed predictions for the level of the unidentified source contribution to the EGRB involve important uncertainties: since no low-energy counterparts have been identified, we have no estimates of distance, and therefore no estimates of the gamma-ray luminosities of these sources. As a result, very few constraints can be placed on their cosmic distribution and evolution. However, very simple estimates can offer some guidance on whether ignoring this EGRB component may be a safe assumption to make. For example, we can use the number of unidentified sources, the minimum flux resolvable by EGRET, and the observed intensity to the extragalactic gamma-ray background to place rough limits on the distance scales associated with resolved and unresolved unidentified sources so that unresolved sources do not overproduce the background. A population of unbeamed, non-evolving, single-luminosity sources uniformly distributed in Euclidian space are resolvable out to a distance $D$ by an instrument of number flux sensitivity $F_{\rm min}$. The relation between $D$, $F_{\rm min}$, and number luminosity $L$ in this case is simply $L=4\pi D^2F_{\rm min}$. If the instrument detects $N$ such sources, their number density $n_{\rm source}$ can be estimated to be $n_{\rm source} = 3N/4\pi D^3$. If the same distribution of sources continues out to a distance $d>D$, the isotropic intensity (photons per unit area per unit time per unit solid angle) from the {\em unresolved} members of this population will be $I_{\rm unres} = \int_D^d dI_{\rm shell}$, where $dI_{\rm shell} = (1/4\pi) (n_{\rm source}4\pi r^2 dr) L/(4\pi r^2)$ is the contribution from sources within a spherical shell located at a distance $r$ from the observer. Substituting our results for $n_{\rm source}$ and $L$ above, and performing the integral, we obtain \begin{equation} I_{\rm unres} = 3NF_{\rm min}(d-D)/4\pi D\,. \end{equation} If we require that the unresolved emission from this population does not exceed the EGRB observed by EGRET ($I_{\rm unres} \leq I_{\rm EGRB}$), we obtain $I_{\rm EGRB} \geq 3NF_{\rm min}(d-D)/4\pi D$. Substituting $N\sim 100$ for the population of extragalactic unidentified sources (see discussion in \S \ref{samples}), $F_{\rm min} \sim 10^{-7} {\rm \, ph \, cm^{2} \, s^{-1}}$ for the sensitivity of EGRET, and $I_{\rm EGRB} \sim 10^{-5} {\rm \, ph \, cm^{2} \, s^{-1} \, sr^{-1}}$ (Sreekumar et al. 1998), we obtain $d \lesssim 6D$. This result implies that the largest distance out to which such a distribution of objects persists cannot be larger than a few times the distance out to which these objects are currently resolved, since in any other case these sources would overproduce the EGRB. It is therefore conceivable that the unidentified source contribution to the EGRB is significant, if not dominant. In this work, we approach the problem from a purely empirical point of view. Instead of attempting to {\em predict} the level of a diffuse component due to unresolved objects of the same class as unidentified EGRET sources, we try to assess whether there are any empirical indications that this component is, in fact, minor. We construct samples of unidentified sources which, based on their sky distribution, are likely to consist mostly of extragalactic objects. Under the assumption that the majority of these sources can be treated as members of a single class of gamma-ray emitters, we seek to answer the following three questions: (1) Is it likely that unresolved objects of the same class could have a significant contribution to the EGRB at least in some energy range? (2) How would the collective spectrum of their emission compare to the measured spectrum of the EGRB deduced from EGRET observations? (3) How are GLAST observations expected to improve our understanding of the nature of unidentified sources, based on the insight gained from our analysis? This paper is structured as follows. In \S \ref{samples} we discuss the samples of resolved unidentified sources used in our analysis. Our formalism for constructing the collective emission spectrum of unresolved unidentified sources is presented in \S \ref{formalism}. Inputs from EGRET data used in our analysis are described in \S \ref{inputs}. In \S \ref{results} we describe our results, and in \S \ref{sec:GLAST} we discuss how these results are expected to improve once GLAST observations become available. Finally, we summarize our conclusions in \S \ref{sec:disc}. | \label{sec:disc} In this work, we have used a purely empirical model to explore the possibility that unresolved gamma-ray sources of the same class as unidentified EGRET sources have an appreciable contribution to the EGRB. We have argued that some unidentified source contribution to the gamma-ray background is guaranteed. We have additionally found that (1) if most high-latitude unidentified sources are assumed to be extragalactic, a one order of magnitude extension of the cumulative flux distribution to lower energies without breaks implies a significant contribution to the EGRB, at least at the lower part of the EGRET energy range; and (2) the spectrum of the cumulative emission of such unresolved sources would be very consistent with the observational determination of Strong et al.\ (2004a) of the EGRET EGRB within systematics. We emphasize that the purpose of this study is not to estimate, even at the order-of-magintude level, the diffuse flux expected from extragalactic unresolved unidentified sources, but rather to place constraints on the flux distribution of these objects under the constraint that the measured background should not be exceeded in any energy interval. Our treatment is therefore different in purpose and spirit from most past work aiming to estimate the level of the contribution of different populations to the extragalactic gamma-ray background. Such work has been traditionally of two types. The first involves population models built from some understanding of the physics of the sources (such as, e.g., the models of M\"{u}cke \& Pohl 2000 and Dermer 2007 in the case of blazars; Miniati 2003, Keshet et al.\ 2003, Blasi, Gabici, \& Brunetti 2007 in the case of clusters of galaxies; Lichti, Bignami \& Paul 1978 and Pavlidou \& Fields 2002 in the case of normal galaxies). The second involves models of the population luminosity function based on our knowledge of the source population from other wavelengths and normalized to fit EGRET data (such studies require a sample of detected, identified members, and are therefore applicable to blazars only, e.g. Chiang et al.\ 1995, Stecker \& Salamon 1996, Mukherjee \& Chiang 1999, Narumoto \& Totani 2006). In our case, lacking any knowledge of the physics of sources as well as of even the bright end of the luminosity function (since in absence of identifications no redshift and hence no luminosity can be derived for any of the sources), we have reversed the problem. Instead of using some assumed cosmic evolution for the sources to derive the expected level of contribution to the EGRB, as in all of the investigations mentioned above, we have used the tightest possible constraint on the allowable EGRB contribution of these sources (the observed EGRET background) to constrain the flux distribution of the unresolved sources. Our conclusions for the overall expected unresolved unidentified source intensity come from the observation that our constraints on the flux distribution are indeed very tight; hence, the contribution of the unidentified sources to the EGRB is likely to be high. Our analysis suggests that any model of the EGRB would be incomplete without some treatment of the unidentified source contribution. The results of our empirical model therefore motivate the pursuit of specific population models for the unidentified sources. Although such models involve a more restrictive set of assumptions and increased uncertainties, they can provide more concrete predictions for the luminosity function of unresolved objects. Once a luminosity function model is assumed, and under the assumption that unresolved members of this class do indeed contribute most of the extragalactic diffuse emission, additional estimates can be made regarding the distance scales associated with this population. The dynamical range in fluxes of the {\em resolved} members of the population, as can be seen in Fig. \ref{lognlogs}, is smaller than an order of magnitude. For a single-luminosity population, this corresponds to a factor of 3 dynamical range in distance, which increases or decreases if the typical source luminosity increases or decreases with increasing redshift respectively. Such arguments can constrain not only the distance scales, but also the number density, intrinsic brightness, and evolution of the resolved and unresolved objects, if in a model-dependent fashion. Additional constraints and predictions, once a luminosity function has been assumed, can be derived through the multi-messenger and multi-wavelength approach. If the emission from extragalactic unidentified sources is primarily of hadronic origin, it will be accompanied by neutrinos at comparable fluxes, and that may have potentially observable consequences in the TeV range for future km$^3$ neutrino detectors such as IceCube and Km3Net. If on the other hand the gamma-ray emission is primarily leptonic, it will be accompanied by X-ray emission, and their extragalactic background flux in the X-ray band may provide an additional constraint. We will pursue such models and calculations in an upcoming publication. In this work we have tried, where possible, to make assumptions, which, if anything, {\em underestimate} the possible contribution of unresolved unidentified sources to the EGRB. An exception to this general trend is our working assumption that the majority of the resolved, high-latitude 3EG unidentified sources belong to a single class. It is conceivable that instead, the resolved unidentified sources are a collection of members of several known and unknown classes of gamma-ray emitters. In this case, it is still likely that the summed contribution of unresolved members of all parent classes to the diffuse background is significant. However, the construction of a single cumulative flux distribution from all sources and its extrapolation to lower fluxes is no longer an indicative test for the importance of such a contribution. Although we have argued that the contribution of unresolved unidentified sources to the EGRB is likely to be important or even dominant, the cumulative emission spectrum derived here and presented in Fig.\ \ref{ress} is only an upper limit, as it was derived demanding that the observed EGRB is not exceeded at any energy above 100 MeV. The contribution of unresolved unidentified sources is further constrained by allowing for the presence of the guaranteed contributions of unresolved normal galaxies and blazars. Finally, it is noteworthy that we have found no evidence for an inconsistency between the population properties of high-latitude unidentified sources and those of blazars. The presence of a yet-unknown population of extragalactic high-energy emitters among the high-latitude EGRET unidentified sources remains one of the most tantalizing possibilities in GeV astronomy and one of the most exciting prospects for the GLAST era. However, the presently available data on the spectral index distribution and the cumulative flux distribution of these sources (this work), as well as their variability properties (e.g., Nolan et al.\ 2003), are consistent with their being members of the blazar population. If indeed a large number of members of the blazar class are present among the entragalactic unidentified sources, then this could have potentially serious effects on our understanding of the redshift distribution of resolved blazars and consequently of the blazar luminosity function. | 7 | 10 | 0710.0619 |
0710 | 0710.0796_arXiv.txt | We present the result of a photometric and Keck-LRIS spectroscopic study of dwarf galaxies in the core of the Perseus Cluster, down to a magnitude of M$_{\rm B}$ $= -12.5$. Spectra were obtained for twenty-three dwarf-galaxy candidates, from which we measure radial velocities and stellar population characteristics from absorption line indices. From radial velocities obtained using these spectra we confirm twelve systems as cluster members, with the remaining eleven as non-members. Using these newly confirmed cluster members, we are able to extend the confirmed colour-magnitude relation for the Perseus Cluster down to M$_{\rm B}$ $= -12.5$. We confirm an increase in the scatter about the colour magnitude relationship below M$_{\rm B}$ $= -15.5$, but reject the hypothesis that very red dwarfs are cluster members. We measure the faint-end slope of the luminosity function between M$_{\rm B}$ $= -18$ and M$_{\rm B}$ $= -12.5$, finding $\alpha$ $= -1.26$ $\pm$ 0.06, which is similar to that of the field. This implies that an overabundance of dwarf galaxies does not exist in the core of the Perseus Cluster. By comparing metal and Balmer absorption line indices with $\alpha$-enhanced single stellar population models, we derive ages and metallicities for these newly confirmed cluster members. We find two distinct dwarf elliptical populations: an old, metal poor population with ages $\sim$ 8 Gyr and metallicities $[{\rm Fe/H}]$ $<$ $-0.33$, and a young, metal rich population with ages $<$ 5 Gyr and metallicities $[{\rm Fe/H}]$ $>$ $-0.33$. Dwarf galaxies in the Perseus Cluster are therefore not a simple homogeneous population, but rather exhibit a range in age and metallicity. | Faint, low mass galaxies are the most numerous galaxy type in the universe, and are thus fundamental in understanding galaxy formation. However, their low luminosities and low surface brightness make detailed studies of them difficult. Because dwarfs are so common, any galaxy evolution/formation theory must clearly be able to predict and describe the properties of these galaxies. The luminosity functions of nearby galaxies in all environments reveal that dwarf galaxies are far more common than brighter galaxies. Within galaxy clusters there furthermore appears to be an overdensity of dwarf galaxies when compared to the field. The origin of these extra dwarfs, or if this excess is even real, remains unknown. The cluster luminosity function itself is not universal, but is strongly dependent on environment \citep{sab03}, with the luminosity function often steeper in the more diffuse outer regions of clusters than in the denser cluster cores. The luminosity function also depends on the individual cluster, with each one having a characteristic luminosity function. For example, the Fornax Cluster is compact and rich in early-type galaxies, with a flat luminosity function faint-end slope of $\alpha = -1.1$ \citep{mie}. The Virgo Cluster has a high abundance of spiral galaxies and has a steeper luminosity function with a faint end slope of $\alpha = -1.6$ (\citealt{trenth02,sab03}). We know that both Local Group (LG) dwarf elliptical and dwarf spheroidal galaxies display varying star formation histories, with metal poor populations as old as the classical halo globular galaxies, but with evidence for star formation as little as 2-3 Gyr ago. Some low-mass LG galaxies such as Sagittarius, contain surprisingly high metal rich stellar populations considering their luminosities, as is also seen in clusters of galaxies. Previous spectroscopic and ground based imaging has revealed that dwarfs in the core of clusters are not a simple homogeneous population, with cluster core galaxies fainter than M$_{\rm B}$ $= -15$ containing a mix of metallicities and ages (\citealt*{Po,RSMP,c1}). The formation scenarios for dwarf galaxies fall into two categories. The first of these is that dwarfs are old, primordial objects (hierarchical model) (e.g. \citealt{wf}). The second is that they have recently evolved or transformed from a progenitor galaxy population (e.g. \citealt*{moore}). Within hierarchical models of galaxy formation, dwarf galaxies are formed from small density fluctuations in the early universe. The lowest mass systems (i.e. dwarf galaxies) form first, and massive galaxies are built from these in mergers. If star formation occurred soon after the gravitational collapse of initial density perturbations, dwarf galaxy halos would be amongst the first objects to form (e.g. \citealt{ds}). Theory also predicts that such systems would form first in the densest environments, although cluster dEs would not necessarily contain the first stars formed in the universe. These halos could have formed their stars quite recently, if star formation was suppressed in some way \citep{tully}. Dwarfs in clusters, on the other hand, could also be the remnants of stripped discs or dwarf irregulars. Through galaxy harassment and interactions, these progenitors become morphologically transformed into dEs. Cluster galaxies become stripped of their interstellar medium, and become dynamically heated by high-speed interactions with other galaxies and the gravitational potential of the cluster. To compensate, the galaxy loses stars, and over time the spiral can morphologically transform into a dE \citep*{c3}. This model is supported by recent observations of embedded discs in dEs in the Virgo and Fornax clusters (\citealt*{Bar,DeR}). Distant clusters at z $\approx$ 0.4 are also filled with small spiral galaxies, but this population is largely absent in nearby clusters, where dwarf spheroidals make up the faint-end of the luminosity function \citep{moore}. Downsizing \citep{DeL} is the observed trend that star formation occurs later, and over more extended timescales, in smaller galaxies. In this scenario, dEs and lower mass galaxies formed or entered clusters after the giant galaxies. Low mass galaxies on average end their star formation after the giant elliptical galaxies. For example, the faint end of the red sequence in clusters is not formed until z $<$ 0.8 \citep{DeL}, implying the luminosity weighted stellar populations of lower mass galaxies are younger than the stars in the giant elliptical galaxies. This seems to contradict the hierarchical method of galaxy formation, where dwarf systems form first. Also, dwarf elliptical galaxies are preferentially found in dense cluster environments, with few examples of isolated dEs, again contradicting the hierarchical model of galaxy formation. However, some dEs in clusters contain old stellar populations, so it is likely that multiple formation methods are needed to explain the origin of cluster dwarfs. The colour-magnitude relation (CMR) for cluster galaxies is well defined at bright magnitudes and forms a tight sequence. However, it is unclear what the shape of the red sequence is at fainter magnitudes \citep{an}. \citet*{c2} find that galaxies at the brightest 4 magnitudes of Perseus obey the CMR, with the fainter candidate members showing significant scatter from the relation. In contrast, \citet{an} and others find no deviation from the CMR at fainter magnitudes in other clusters. Instead, the faint red sequence is an extrapolation of what is observed at brighter magnitudes. With spectroscopy, we can test this directly by finding confirmed dwarf members of the cluster to establish the true form of the colour magnitude relation. Spectroscopy can also be used to measure the faint-end of the luminosity function in galaxy clusters, which is fundamental in describing the galaxy population. The luminosity function contains important information on the formation and evolution of galaxies, and at low redshifts contains the combined influence of the galaxy initial mass function, and the effects of any evolutionary processes that have taken place in the cluster since its formation. In this paper, we use deep Keck spectroscopy to determine cluster membership for dwarfs at M$_{\rm B}$ $<$ $-16$, and present the ages and metallicities of these galaxies based on the strengths of Balmer and metal absorption lines in their spectra. We confirm that twelve dwarfs are Perseus Cluster members, while the remaining objects are background galaxies. Using these results, we examine the colour-magnitude relation for the Perseus Cluster down to M$_{\rm B}$ $= -12.5$. We also measure the luminosity function faint-end slope $\alpha$ based on our cluster membership results. Our main conclusion is that some dwarf galaxies in Perseus have old, metal poor populations, whilst others are younger, metal rich systems. This suggests that dwarf galaxies in the core of the Perseus Cluster are not a simple, homogeneous population, but require multiple scenarios for their formation. This paper is organised as follows: in $\S$ 2 we discuss the observations and the selection criteria for the dwarf galaxies, $\S$ 3.1 identifies the cluster members, $\S$ 3.2 presents the colour magnitude relation, and the central luminosity function is presented in $\S$ 3.3. In $\S$ 3.4 we derive luminosity weighted ages and metallicities for the newly confirmed cluster members. A discussion of the main results is presented in $\S$ 4 and these results are summarised in $\S$ 5. We assume the distance to the Perseus Cluster is 77 Mpc throughout this paper. | We have analysed Keck/LRIS spectra for twenty-three dwarf galaxy candidates in the Perseus Cluster, from which we have confirmed cluster membership for twelve systems based on radial velocities measured from absorption lines. We extend the confirmed member colour-magnitude relation for Perseus down to M$_{\rm B}$ $= -12.5$, finding that the slope of the colour-magnitude relation becomes bluer when the low-surface brightness dwarf galaxies are included. The fainter dwarfs also scatter more from the colour-magnitude relation, following the trend observed by \citet{RS} for low-mass galaxies in Fornax and Coma. This scatter can be interpreted as a spread in the metallicity distributions of dwarf galaxies, which has been inferred for dwarf galaxies in other clusters by \citet{RSMP} and \citet{Po}. After removing non-members from the $B$-band luminosity function for the Perseus Cluster we find a faint-end slope $\alpha$ $= -1.26$ $\pm$ 0.06, similar to the field. Previous studies of other galaxy clusters have found that the dwarf-to-giant ratio is a function of local projected galaxy density. Extending this study to the outer regions of Perseus would enable us to see how the luminosity function changes with the local galaxy density. Colours, morphologies, and central surface brightnesses are not sufficient criteria to confirm cluster membership, as this work has shown. Cluster membership cannot be confirmed without spectroscopy, so the faint-end luminosity functions calculated for galaxy clusters where membership has not been confirmed spectroscopically are most likely artificially steepened by non-members resulting in higher dwarf-to-giant ratios. By comparing the observed dwarf spectral absorption indices with population synthesis models of \citet{tho}, we derive luminosity-weighted ages and metallicities for the dwarf galaxies. A range of ages is observed, ranging from older than 8 Gyr to younger than 5 Gyr. The metallicity distribution of the faint cluster members is also not that of a simple homogeneous population, with the younger galaxies typically having higher metallicities. More observations of dwarf galaxies in rich, nearby cluster environments are required in order to help improve our understanding of the formation scenarios for cluster dwarfs. Further spectroscopic work with a larger sample would also allow better constraints on the ages and metallicities of cluster dwarf galaxies and would help with modelling the formation of such galaxies. | 7 | 10 | 0710.0796 |
0710 | 0710.2793_arXiv.txt | SGRs are galactic X-ray stars that emit, during sporadic times of high activity, a large number of short-duration (around 0.1 s) bursts of hard X-rays (Duncan and Thompson, 1992). A SGR is thought to be a magnetar, being a strongly magnetized neutron star powered by a very strong magnetic field ($\ge$ 10$^{15}$ Gauss). On 27 December 2004 a powerful burst of X- and $\gamma$-rays from one of the most highly magnetized neutron stars (SGR 1806-20) of our Galaxy reached the Earth's environment (Hurley et al., 2005). The Solar system received a shock, which is thought to be due to a cataclysm in the magnetar that caused it to emit as much energy in two-tenths of a second as the Sun gives off in 250,000 years. The signature of this event on the Earth's magnetic field has not previously been investigated. Here, we present the first results of the magnetar footprints on magnetic data recorded by near-Earth satellites. The magnetar SGR 1806-20 is the third such event ever recorded along with two others that were noted in 1979 and 1998 (Mazets et al., 1979; Hurley et al., 1999). Several properties of this magnetar flare are relevant to our study. Firstly, a precursor of $\sim$ 1 s was observed 142 s before the flare, with a roughly flat-topped profile (Hurley et al., 2005). The intensity of the main initial spike saturated all X- and $\gamma$-ray detectors. However, particle detectors on board of RHESSI and Wind spacecraft (Boggs et al., 2004; Mazets et al., 2004) were able to record reliable measurements. Several instruments designed for other purposes provided important information, as Geotail (Terasawa et al., 2005) and Cluster/Double star (Schwartz et al., 2005). The first spike was followed by a tail lasting 380 s, during which 7.56 s pulsations were clearly observed, by the $\gamma$-ray detectors on board of RHESSI (Hurley et al., 2005). Secondly, a disturbance of the Earth's ionosphere was simultaneously observed with the detection of the burst from SGR 1806-20 (Inan et al., 2005). This sudden ionospheric disturbance (SID) was recorded as a change in the signal strength from very low frequency (VLF) radio transmitters, being noticed by stations around the globe (Campbell et al., 2005). These changes in the radio signal strength were caused by X-rays arriving from SGR 1806-20, which ionized the upper atmosphere and modified the radio propagation properties of the Earth's ionosphere (see clearing house of SID data associated with SGR 1806-20 flare at http://www.aavso.org/observing/programs /solar/sid-sgr1806.shtml). One such observation of this ionospheric signature resides within a 21.4 kHz signal that originates in Hawaii and propagates along an ionosphere wave guide to Palmer Station, Antarctica (Inan et al., 2005). This wave guide is some $\sim$ 10,000 km in path length (Inan et al., 2005). As explained above, this is not a direct radio detection of SGR 1806-20 (see also http://gcn.gsfc.nasa.gov/gcn3/2932.gcn3). Moreover, due to the sub-burst longitude and latitude (Inan et al., 2005) and to the geographical distribution of LF/VLF beacons and monitoring stations, this burst was not detected by active monitoring stations in Germany, Australia, or Canada (Campbell et al., 2005). Here, we note that ionospheric disturbances were also reported in the case of the magnetar observed in 1998 (Inan et al., 1999). In the case of the 1998 magnetar the flare illuminated the nightside of the Earth and ionized the lower ionosphere to levels usually found only during daytime. The magnetar responsible for the 2004 burst was about the same distance as the magnetar responsible for the 1998 burst, but within 5.25$^\circ$ of the Sun as viewed from Earth. Therefore its $\gamma$-rays arrived on the dayside of our planet. The 2004 flare changed the ionic density at an altitude of 60 km by six orders of magnitude (Inan, 2006). It is thus plausible that this change in the ionospheric conductivity can cause oscillating perturbations in the current-generated magnetic field. The thrust of this study is to find signatures associated with the explosion of the magnetar SGR 1806-20 within satellite measurements of Earth's electromagnetic field. Currently, the Earth's electromagnetic field is monitored by a number of Low Earth Orbit (LEO) satellite missions. After the launch of \O rsted satellite in 1999, the knowledge of the near-Earth electromagnetic field has been dramatically improved (Hulot et al., 2002; L\"uhr et al., 2002; Maus et al., 2002; Tyler et al., 2003; Balasis et al, 2004). Since 2000, \O rsted, CHAMP and SAC-C satellites have offered a continuous flow of high quality magnetic field measurements. Additionally, the DEMETER satellite provides 1 Hz energetic electron detector data. Finally, let us note that all these LEO magnetic missions are flying between the Earth's surface, where the temporal variations of the magnetic field are continuously monitored by geomagnetic observatories, and the magnetosphere, where an in-situ investigation of the three-dimensional and time-varying phenomena is done by the four identical spacecraft of Cluster II mission. | The effect of the SGR1806-20 flare on the Earth's magnetic field was not large, but it was detectable. This first attempt to find a magnetar signature in the geomagnetic field clearly indicates that the high resolution CHAMP magnetic data are optimal to capture the extremely bright flare from SGR 1806-20. Indeed, during the first half of the decay phase of the flare a 7.5 s periodicity is observed in the magnetic field over a magnetically quiet period, near the South Pole at 400 km altitude. This observation can be explained by a mechanism through which the oscillating flux of ionizing $\gamma$-rays could alter the ionospheric conductivity and hence cause oscillating perturbations in the current-generated magnetic field. An attempt to verify this hypothesis for the two previously recorded giant flares was not possible since no magnetic satellite missions were operating in LEO at that time (i.e., in March 1979 and August 1998). Of course, there are many spacecraft carrying magnetometers within the Solar system, but very few near planetary ionosphere. For example, the wavelet analysis was performed on magnetometer data from the Cluster II mission that probes the Earth's magnetosphere. In order to be able to visualize, in the wavelet power spectrum graph, any significant disturbances of the magnetic field, the power spectral density of the signal was amplified by a factor of $2^6$ (in comparison to the corresponding spectral density values of the CHAMP data). Although there are some indications for a weak pulsation-like signal at $\sim$ 8 s, the fact that this signal is almost two orders of magnitudes weaker than the one observed in CHAMP data favors the hypothesis of an ionospheric origin for the signature found in CHAMP data. Furthermore, our analysis can be extended to 1 Hz magnetic data provided by ground-based magnetic observatories, but only a small number of them provide such high resolution sampling nowadays. Data provided by 12 Canadian observatories, for which 1 Hz values are available over the period we are interested in, were also analyzed. For five of these observatories, missing data or high-level noise, made it difficult to apply the wavelet technique. For the others, no conclusive evidence for a signature related to SGR 1806-20 exists. Analyzing other magnetic data with such a powerful tool as wavelets techniques could be relevant for understanding the impact that giant flares have on the terrestrial and other planetary magnetic fields. However, the main difficulty in such studies is due to the availability and quality of magnetic data. For instance, wavelet analysis of Mars Global Surveyor mission magnetic measurements on 27/12/04 was not able to detect any of the magnetar features due to the inadequate sampling rate: only 3 s data are now available (Michael Purucker, pers. comm. 2005). | 7 | 10 | 0710.2793 |
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0710 | 0710.3795_arXiv.txt | Massive black hole binary coalescences are prime targets for space-based gravitational wave (GW) observatories such as {\it LISA}. GW measurements can localize the position of a coalescing binary on the sky to an ellipse with a major axis of a few tens of arcminutes to a few degrees, depending on source redshift, and a minor axis which is $2 - 4$ times smaller. Neglecting weak gravitational lensing, the GWs would also determine the source's luminosity distance to better than percent accuracy for close sources, degrading to several percent for more distant sources. Weak lensing cannot, in fact, be neglected and is expected to limit the accuracy with which distances can be fixed to errors no less than a few percent. Assuming a well-measured cosmology, the source's redshift could be inferred with similar accuracy. GWs alone can thus pinpoint a binary to a three-dimensional ``pixel'' which can help guide searches for the hosts of these events. We examine the time evolution of this pixel, studying it at merger and at several intervals before merger. One day before merger, the major axis of the error ellipse is typically larger than its final value by a factor of $\sim 1.5-6$. The minor axis is larger by a factor of $\sim 2-9$, and, neglecting lensing, the error in the luminosity distance is larger by a factor of $\sim 1.5-7$. This large change over a short period of time is due to spin-induced precession, which is strongest in the final days before merger. The evolution is slower as we go back further in time. For $z = 1$, we find that GWs will localize a coalescing binary to within $\sim 10\ \mathrm{deg}^2$ as early as a month prior to merger and determine distance (and hence redshift) to several percent. | \label{sec:intro} Among the most important sources of gravitational waves (GWs) in the low-frequency band of space-based detectors are the coalescences of massive black hole binaries (MBHBs). Binaries containing black holes with masses in the range $10^4 - 10^7\,M_\odot$ are predicted to form through the hierarchical growth of structure as dark matter halos (and the galaxies they host) repeatedly merge; see, for example, {\citet{svh07}} and {\citet{mhsa07}} for recent discussion. The {\it Laser Interferometer Space Antenna} ({\it LISA}) is being designed to have a sensitivity that would allow detailed measurement of the waves from these binaries. ``Intrinsic'' parameters --- the masses and spins of the black holes which compose the binary --- should be determined with very high accuracy, with relative errors typically $\sim 10^{-3}$ to $10^{-1}$, depending on system mass and redshift; see Lang \& Hughes (2006, hereafter Paper I) for recent discussion. By measuring an ensemble of coalescences over a range of redshifts, MBHB GWs may serve as a kind of structure tracer, tracking the growth and spin evolution of black holes over cosmic time. ``Extrinsic'' system parameters, describing a binary's location and orientation relative to the detector, are also determined by measuring its GWs. In Paper I, we showed that a binary's position on the sky can be localized at $z = 1$ to an ellipse with a major axis of a few tens of arcminutes and a minor axis a factor of $2-4$ smaller. At higher redshift ($z = 3-5$), these values degrade by a factor of a few, reaching a few degrees in the long direction and tens of arcminutes to a degree or two in the short one. We also found that, neglecting errors due to weak gravitational lensing, a source's luminosity distance typically can be determined to better $1\%$ at low redshift ($z = 1$), degrading to several percent at higher redshift ($z = 3-5$). The intrinsic ability of GWs to determine the distance to a coalescing binary is phenomenal. Coalescing MBHB systems constitute exquisitely well calibrated distance measures, with the calibration provided by general relativity. Unfortunately, this percent-level or better accuracy could only be achieved if we measured MBHB coalescences in an empty universe. In our universe, weak lensing will magnify or demagnify the GWs, and we will infer a luminosity distance smaller (for magnification) or larger (for demagnification) than the true value. This phenomenon affects all high-redshift standard candles. Its impact on Type Ia supernovae in particular has been discussed in detail {\citep{frieman97,hw98,h98}}. It will not be possible to correct for this effect {\citep{dhcf03}} since much of the lensing ``noise'' arises from structure on subarcminute scales that is not probed by shear maps (which map the distribution of matter on scales greater than $1^\prime$ or so). Since we will not know the extent of the magnification when we measure MBHB waves, we must simply accept the fact that lensing introduces a dispersion of several percent in determining the distance to these GW events (see, e.g., Wang et al. [2002] to compute this dispersion as a function of redshift). When we quote distance measurement errors, we will typically quote only the intrinsic GW measurement error, neglecting lensing's impact. When the intrinsic GW distance error is $\lesssim 5\%$, lensing will blur it to the several percent level. Note that a source's redshift $z$ {\it cannot} be directly determined using only GWs. Gravitational wave measurements infer system parameters through their impact on certain dynamical timescales, such as orbital frequencies and the rate at which these frequencies evolve. Since these time scales all suffer cosmological redshift, $z$ is degenerate with other parameters. For example, any mass parameter $m$ is actually measured as $(1 + z)m$. However, if the binary's luminosity distance is determined, its redshift can then be inferred by assuming a cosmography. For most binaries, the redshift can be determined to several percent (with an error budget dominated by gravitational lensing\footnote{At redshifts $z \lesssim 0.3$, the error is actually dominated by peculiar velocity effects {\citep{kfhm06}}; however, the event rate is probably negligible at such low redshifts. As such, we will focus on gravitational lensing as the main source of systematic redshift error.}). We thus expect that GW measurements will locate a binary to within a three-dimensional ``GW pixel'' which at $z = 1$ has a cross-sectional area of $\sim 10^{-2}$ to $10^{-1}\ \mathrm{deg}^2$ and a depth $\Delta z/z \sim $ several percent. It is anticipated that there will be great interest in searching the GW pixel for electromagnetic (e.g., optical, X-ray, radio) counterparts to MBHB GW events. Finding such a counterpart would be much easier if galactic activity were catalyzed in association with the coalescence {\citep{kfhm06}}. The nature of that activity is likely to depend rather strongly on the mass of the coalescing system \citep{dssch06}. For example, {\citet{an02}} predict strong outflows and galactic activity prior to the final black hole merger as the smaller member of the binary drives gas onto the larger member, consistent with the high-mass ($M_{\rm tot} \gtrsim 10^7\,M_\odot$) predictions of \citet{dssch06}. {\citet{mp05}} describe an X-ray afterglow that would ignite after gas refills the volume that is swept clean by the coalescing binary; Dotti et al.\ predict this outcome for smaller systems ($M_{\rm tot} \lesssim \mbox{several}\times 10^6\,M_\odot$). Recent work by {\citet{bp07}} suggests that the final burst of radiation from a coalescing binary (which can convert $\sim 10\%$ of the system's mass to GWs very suddenly) may excite radial waves, and consequently electromagnetic variability, in an accretion disk due to the quick change in the disk's Keplerian potential. Such a signature may be essentially mass independent. On the other hand, the coalescence may be electromagnetically quiet, in which case we face the potentially daunting task of searching the three-dimensional pixel for galaxies with morphology consistent with a (relatively) recent merger, or that have a central velocity dispersion $\sigma$ consistent with the inferred final black hole mass (assuming that the $M_{\rm BH}-\sigma$ relation [Ferrarese \& Merritt 2000; Gebhardt et al. 2000] holds at the redshift of these sources, and so soon after merger). If the host galaxy or some other counterpart can be identified, we could then contemplate combining GW information with electromagnetic data. For instance, combining {\it LISA} mass measurements with the luminosity of the counterpart may allow us to directly measure the Eddington ratio $L/L_{\mathrm{Edd}}$ {\citep{kfhm06}}. Identifying a counterpart would also allow us to more accurately characterize the system. For example, much of the intrinsic luminosity distance error is due to correlations between distance and sky position. Finding an electromagnetic counterpart essentially determines a binary's position precisely, breaking those correlations. Previous studies have found that intrinsic distance error can be reduced by almost an order of magnitude if the position is known {\citep{h02,hh05}}. (Lensing errors still dominate in such a case, so that the distance remains determined only at the few percent level.) A counterpart may also make it possible to simultaneously determine a source's luminosity distance and redshift. Such a ``standard siren'' (the GW analog of a standard candle) would very usefully complement other high-redshift standard candles {\citep{hh05}}, such as Type Ia supernovae {\citep{p93, rpk95, wgap03}}. A direct measurement of redshift will also break the mass-redshift degeneracy more accurately than can be done with just the luminosity distance and some assumed cosmological parameters. Breaking this degeneracy is critical when studying the growth of black holes with cosmic time {\citep{h02}}. Many analyses {\citep{c98,h02,v04,bbw05,hh05}} have quantified how well {\it LISA} can determine MBHB parameters, including sky position and distance, using maximum likelihood Fisher matrix estimation. Our results from Paper I, given earlier, include the effects of ``spin-induced precession'' --- precession of both the orbital plane and the individual spins of the black holes due to post-Newtonian spin-interaction effects. A significant result from that analysis is that spin-induced precession improves sky position accuracy by about half an order of magnitude in each direction versus previous analyses. This result can be partially understood as due to the breaking of a degeneracy between position and orientation angles: thanks to precession, the binary's orientation evolves with time and can be untangled from sky position. This effect was already known due to pioneering work by {\citet{v04}}; by taking the analysis to higher order and considering a broader range of sources, we were able to show that this improvement held for essentially all astrophysically interesting MBHB sources. The purpose of this paper is to examine the localization of MBHB systems more thoroughly, in particular how the GW pixel evolves as the final merger is approached. Paper I only presented results for measurements that proceed all the way to merger. It will clearly be of some interest to monitor potential hosts for the binary event some time before the merger happens; if nothing else, telescopes will need prior warning to schedule observing campaigns. Understanding the rate at which localization evolves can also have an important impact on the design of the {\it LISA} mission, clarifying how often it will be necessary to downlink data about MBHB systems in order to effectively guide surveys. Our main goal is to understand for what range of masses and redshifts prior localization of a binary using GWs will be possible. A previous analysis by Kocsis et al.\ (2007b; hereafter K07) also examined this problem in great detail, but without including the impact of spin-induced precession. One of our goals is to see to what extent precession physics changes the conclusions of K07. We find that precession has a fairly small impact on the time evolution of the GW pixel except in the last few days before the final merger, at which point its impact can be tremendous. Precession typically changes the area of the sky position error ellipse by a factor of $\sim 3-10$ (up to $\sim 60$ in extreme cases) in just the final day. This is in accord with the predictions of K07 (and even earlier predictions by N. Cornish 2005, unpublished). The structure of this paper is as follows. First, in \S\ {\ref{sec:background}}, we briefly review the basics of the MBHB gravitational waveform and the parameter estimation formalism that we use; this section is essentially a synopsis of relevant material from Paper I. Section {\ref{sec:intrinsicGW}} reviews the form of the GWs that we use in our analysis, while \S\ {\ref{sec:extrinsicGW}} describes how those waves are measured by the {\it LISA} constellation. We describe the measurement formalism we use in \S\ {\ref{sec:formalism}}. In \S\ {\ref{sec:review}}, we summarize our results from Paper I regarding the final localization accuracy that {\it LISA} can expect to achieve. We turn to a detailed discussion of the time evolution of the GW pixel in \S\ {\ref{sec:results}}. We begin by summarizing the key ideas behind the ``harmonic mode decomposition'' of K07 in \S\ {\ref{sec:khmf}}. This technique cleverly allows calculation of the GW pixel and its time evolution with much less computational effort than our method (albeit without including the impact of spin precession). Unfortunately, we have discovered that some of the approximations used by K07 introduce a systematic underestimate of the final sky position error by a factor of $2 - 4$ or more in angle; the approximations are much more reliable a week or more prior to the black holes' final merger. Modulo this underestimate, the K07 results agree well with a version of our code which does not include spin precession (particularly a week or more in advance of merger, when their underestimate is not severe). K07 thus serves as a useful point of comparison to establish the impact of precession on source localization. Section {\ref{sec:timeevolve}} is dedicated to our results, including comparison to K07 when appropriate. We find that all relevant parameter errors decrease slowly with time until the last day before merger, when they drop more dramatically. This sudden drop is not found in K07, nor is it present in a variant of our analysis that ignores spin precession. It clearly can be attributed to the impact of precession on the waveform. Before this last day, the major axis is $\sim 1.5-6$ times, the minor axis $\sim 2-9$ times, and the intrinsic error in the luminosity distance $D_L$ $\sim 1.5-7$ times bigger than at merger for most binaries (i.e., all except the highest masses). Going back to one week (one month) before merger, these numbers change to $2-9$ ($4-11$) for the major axis, $3-14$ ($5-24$) for the minor axis, and $3-14$ ($5-18$) for the error in the luminosity distance. As a result, for $z = 1$, most binaries can be located within a few square degrees a week before merger and $10\ \mathrm{deg}^2$ a month before merger. The intrinsic distance errors are also small enough this early that $\Delta z/z$ remains dominated by gravitational lensing errors of several percent. Advanced localization of MBHB coalescences thus seems plausible for these binaries; the situation is less promising for sources at higher redshift. As a corollary to our study of the time evolution, we also examine the sky position dependence of errors (\S\ {\ref{sec:angdependence}}). The errors depend strongly on the polar angle with respect to the ecliptic, increasing in the ecliptic plane to as much as $35\%$ over the median for the major axis, $85\%$ over the median for the minor axis, and $15\%$ over the median for errors in the luminosity distance. The errors have a much weaker dependence on the azimuthal angle. When we convert to Galactic coordinates, we find that the best localization regions appear to lie fairly far out of the Galactic plane, offering hope that searches for counterparts will not be too badly impacted by foreground contamination. We conclude this paper in \S\ {\ref{sec:disc}}. Besides summarizing our results, we discuss shortcomings of this analysis and future work which could help to better understand how well GWs can localize MBHB sources. Throughout the paper we set $G = c = 1$; a convenient conversion factor in this system is $10^6 M_\odot = 4.92$ s. When discussing results, we always quote masses as they would be measured in the rest frame of the source. These masses must be multiplied by a factor of $1 + z$ when used in any of the equations describing the system's dynamics or its GWs (particularly the equations of \S\ {\ref{sec:background}}). We convert between distance and redshift using a spatially flat cosmology with $\Omega_\Lambda = 0.75$ (and hence $\Omega_m = 0.25$) and Hubble constant $H_0 = 75\ {\rm km}\ {\rm s}^{-1}\ {\rm Mpc}^{-1}$. | \label{sec:disc} As discussed at length in Paper I, accounting for the general relativistic precession of the angular momentum vectors in an MBHB system has a dramatic impact on what we can learn by observing the system's gravitational waves. Spin-induced precession breaks degeneracies among different parameters, making it possible to measure them more accurately than they could be determined if precession were not present. This has a particularly important impact on our ability to locate such a binary on the sky and to determine its luminosity distance --- the degeneracy between sky angles, distance, and orientation angles is severe in the absence of precession. Our analysis shows that the improvement that precession imparts to measurement accumulates fairly slowly. In using one code which includes the impact of spin precession and a second which neglects this effect, we find little difference in the accuracy with which GWs determine sky position and distance for times more than a few days in advance of the final merger. The difference between the two codes grows quite rapidly in these final days. In the last day alone, the localization ellipse area decreases by a factor of $\sim 3-10$ (up to $\sim 60$ in a few low-mass systems) when precession effects are included. Distance determination is likewise improved by factors of $\sim 1.5-7$ in that final day. Not all of the precession effects occur in the final days. We saw in Figure \ref{fig:axes_evol} that the long tail of small minor axes can be seen, to some degree, throughout the inspiral. We could get lucky and find a binary with a very small value of $2b$ weeks before merger. But the improvement in the median that we found in Paper I appears to take effect only in the final days of inspiral. Therefore, while precession may in fact help improve the {\it final} localization of a coalescing binary by a factor of $\sim 2-10$ in each direction, it will not be much help in {\it advanced} localization of a typical binary. Nevertheless, the pixel sizes that we find are small enough that future surveys should not have too much trouble searching the region identified by GWs, at least over certain ranges of mass and redshift. At $z = 1$, the GW localization ellipse is $\sim 10\ \mathrm{deg}^2$ or smaller for most binaries as early as a month in advance of merger. (At high masses, the ellipse can be substantially larger than this a month before merger, but it shrinks rapidly, reaching a comparable size $1-2$ weeks before merger.) This bodes well for future surveys with large fields of view that are likely to search the GW pixel for counterparts. In addition, GWs determine the source luminosity distance so well that the distance errors we find are essentially irrelevant --- gravitational lensing will dominate the distance error budget for all but the highest masses. As redshift increases, the GW pixel rapidly degrades, particularly for the largest masses. Let us adopt $10\ \mathrm{deg}^2$ (the approximate LSST field of view) as a benchmark localization for which counterpart searches may be contemplated. At $z = 3$, this benchmark is reached at merger for almost the entire range of masses we considered. As little as a day in advance of merger, however, some of the least massive and most massive systems are out of this regime. One week prior to merger, the most massive systems are barely located at all (ellipses hundreds of square degrees or larger). The intermediate masses do best, but even in their cases the positions are determined with $\sim 10\ \mathrm{deg}^2$ accuracy no earlier than a few days in advance of merger. The resolution degrades further at higher redshift. At $z = 5$, systems with $M \gtrsim 6 \times 10^6\,M_\odot$ are not located more accurately than $\sim 30\ \mathrm{deg}^2$ even at merger. Smaller systems are located within $\sim 10\ \mathrm{deg}^2$ at merger, but very few are at this accuracy even one day in advance of merger. The luminosity distance errors also increase, so much that they exceed lensing errors a few days to a week before merger at $z = 3$, and only a day before merger at $z = 5$. This degradation hurts the ability to search for counterparts by redshift and subsequently use them as standard candles. Our main conclusion is that future surveys are likely to have good advanced knowledge (a few days to one month) of the location of MBHB coalescences at low redshift ($z \sim 1 - 3$), but only a day's notice at most at higher redshift ($z \sim 5$). This conclusion may be excessively pessimistic. As mentioned earlier, recent work examining the importance of subleading harmonics of MBHB GWs is finding that including harmonics beyond the leading quadrupole has an important effect on the final accuracy of position determination {\citep{arunetal,ts08}}. For most masses, these analyses show a factor of a few improvement in position, comparable to the improvement that we find when spin precession is added to the waveform model. For high-mass systems, the higher harmonics increase the (previously small) overlap with the {\it LISA} band; consequently, the improvement can be much larger, up to 2 or 3 orders of magnitude in area. Since these two improvements arise from very different physical effects, it is likely that their separate improvements can be combined for an overall improvement significantly better than each effect on its own. We plan to test this in future work (which is just now getting underway). Finally, we have also studied the sky position dependence of {\it LISA}'s ability to localize sources. We have found that the regions of best localization lie fairly far out of the Galactic plane. However, as emphasized by N. Cornish (2007, private communication), a proper anisotropic confusion background might impact this dependence. In our calculations, we have assumed an isotropic background, neglecting the likely spatial distribution of Galactic binaries. Properly accounting for this background is likely to strengthen our conclusion that LISA's ability to ``see'' is best for MBHB sources out of the Galactic plane. | 7 | 10 | 0710.3795 |
0710 | 0710.3106_arXiv.txt | Development of the Aarhus adiabatic pulsation code started around 1978. Although the main features have been stable for more than a decade, development of the code is continuing, concerning numerical properties and output. The code has been provided as a generally available package and has seen substantial use at a number of installations. Further development of the package, including bringing the documentation closer to being up to date, is planned as part of the HELAS Coordination Action. | The goal of the development of the code was to have a simple and efficient tool for the computation of adiabatic oscillation frequencies and eigenfunctions for general stellar models, emphasizing also the accuracy of the results. Not surprisingly, given the long development period, the simplicity is now less evident. However, the code offers considerable flexibility in the choice of integration method as well as ability to determine all frequencies of a given model, in a given range of degree and frequency. The choice of variables describing the equilibrium model and oscillations was to a large extent inspired by \citet{Dziemb1971}. As discussed in Section~\ref{sec:eqmodel} the equilibrium model is defined in terms of a minimal set of dimensionless variables, as well as by mass and radius of the model. Fairly extensive documentation of the code, on which the present paper in part is based, is provided with the distribution package% \footnote{The package is available at \\ {\tt http://astro.phys.au.dk/$\sim$jcd/adipack.n}}. \citet{Christ1991} provided an extensive review of adiabatic stellar oscillations, emphasizing applications to helioseismology, and discussed many aspects and tests of the Aarhus package, whereas \citet{Christ1994} carried out careful tests and comparisons of results on polytropic models; this includes extensive tables of frequencies which can be used for comparison with other codes. | 7 | 10 | 0710.3106 |
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0710 | 0710.2954_arXiv.txt | We have compiled the emission-line fluxes of \ion{O}{1} $\lambda$8446, \ion{O}{1} $\lambda$11287, and the near-infrared (IR) \ion{Ca}{2} triplet ($\lambda$8579) observed in 11 quasars. These lines are considered to emerge from the same gas as do the \ion{Fe}{2} lines in the low-ionized portion of the broad emission line region (BELR). The compiled quasars are distributed over wide ranges of redshift (0.06 $\le z \le$ 1.08) and of luminosity ($-29.8 \le M_{B} \le -22.1$), thus providing a useful sample to investigate the line-emitting gas properties in various quasar environments. The measured line strengths and velocities, as functions of the quasar properties, are analyzed using photoionization model calculations. % We found that the flux ratio between the \ion{Ca}{2} triplet and \ion{O}{1} $\lambda$8446 is hardly dependent on the redshift or luminosity, indicating similar gas densities in the emission region from quasar to quasar. On the other hand, a scatter of the \ion{O}{1} $\lambda$11287/$\lambda$8446 ratios appears to imply the diversity of the ionization parameter. These facts invoke a picture of the line-emitting gases in quasars that have similar densities and are located at regions exposed to various ionizing radiation fluxes. The observed \ion{O}{1} line widths are found to be remarkably similar over more than 3 orders of magnitude in luminosity, which indicates a kinematically determined location of the emission region and is in clear contrast to the case of \ion{H}{1} lines. We also argue about the dust presence in the emission region since the region is suggested to be located near the dust sublimation point at the outer edge of the BELR. | Active galactic nuclei (AGNs) are known to have strong emission lines of various ion species. Among them, the \ion{Fe}{2} emission lines are one of the most prominent features in the ultraviolet (UV) to optical spectrum of many AGNs. They have long been hoped to provide significant information about some aspects of the AGNs and their host environments, e.g., energy budget of the line emission region % and the epoch of the first star formation in the host galaxies. The determination of the first star formation epoch is based on the standard theory that the iron enrichment in galaxies is delayed compared to that of the $\alpha$-elements, such as magnesium, due to their different origins; Type Ia supernovae for iron and Type II supernovae for the $\alpha$-elements \citep{hamann93,yoshii98}. The delay corresponds to the difference in life-times of the progenitors of the two types of supernovae, and is estimated to be 0.3 -- 1 Gyr depending on the host galaxy environments \citep{yoshii96, matteucci01}. Many observations have been devoted to the measurement of \ion{Fe}{2}/\ion{Mg}{2} line flux ratios in high-redshift quasars for this purpose over the last decade \citep[e.g.,][]{elston94,kawara96,dietrich02,dietrich03,iwamuro02,iwamuro04,freudling03,maiolino03}. However, the observed \ion{Fe}{2}/\ion{Mg}{2} ratios show a large scatter, preventing a detection of any significant trend in the Fe abundance as a function of redshift. While a part of the scatter might be due to the difference in the intrinsic Fe/Mg abundance ratio, it is presumable that the diversity of the physical condition within the line-forming gas, affecting line emissivities, is the main cause \citep{verner03, baldwin04}. In the same sense, a change of the observed \ion{Fe}{2}/\ion{Mg}{2} ratio as a function of redshift, if found, should be carefully examined to tell whether it reflects the abundance evolution or the systematic variation of the line emissivity. Thus establishment of a method to probe the line-emitting gas and estimate its physical parameters such as density and incident-ionizing radiation flux has been much awaited. Unfortunately, the Fe$^+$ atom is characterized by an enormous numbers of possible electronic transitions, yielding the ``\ion{Fe}{2} pseudocontinuum'' often observed in AGN spectra, which makes analysis of the observations extremely difficult from both the observational and theoretical viewpoints \citep[e.g., ][hereafter T06]{tsuzuki06}. On the other hand, a promising approach is to use the emission lines emitted by simple atoms in the same region as the \ion{Fe}{2} lines. The most potent lines are \ion{O}{1} and \ion{Ca}{2}, whose co-spatial emergence with \ion{Fe}{2} is indicated by a resemblance of their profiles \citep{rodriguez02a} and by a correlation between the line strengths \citep{persson88}. Note that it is a natural consequence of similar ionization potentials of the relevant ions, i.e., 16.2 eV for \ion{Fe}{2}, 13.6 eV for \ion{O}{1}, and 11.9 eV for \ion{Ca}{2}. The first extensive study of the physical properties of \ion{O}{1} emitting gases in AGNs was presented by \citet{grandi80}, who observed the strongest \ion{O}{1} line, $\lambda$8446, as well as other weaker \ion{O}{1} lines in Seyfert 1 (Sy1) galaxies. He found that \ion{O}{1} $\lambda$8446 lacks the narrow component that characterizes other permitted lines, and concluded that the line is purely a BELR phenomenon. He also suggested that \ion{O}{1} $\lambda$8446 is produced by Ly$\beta$ fluorescence, which was later confirmed by the observation of I Zw 1, the prototype narrow-line Seyfert 1 (NLS1), by \citet{rudy89}. \citet{rodriguez02b} compiled the UV and near-IR \ion{O}{1} lines, namely, $\lambda$1304, $\lambda$8446, and $\lambda$11287, in normal Sy1s and NLS1s in order to investigate their flux ratios. They found that there must be an additional excitation mechanism for \ion{O}{1} $\lambda$8446---besides Ly$\beta$ fluorescence---which they concluded is collisional excitation. As for the \ion{Ca}{2} lines, extensive studies of Sy1s were presented by \citet{persson88} and \citet{ferland89}. We have observed seven quasars at redshifts up to $\sim$1.0 for the purpose of obtaining the UV and near-IR \ion{O}{1} and \ion{Ca}{2} lines, thus extending the previous studies to include quasars at high redshifts. Photoionization model calculations were performed and compared with the observations, which led us to conclude that the \ion{O}{1} and \ion{Ca}{2} lines are formed in a gas with density $n_{\rm H} = 10^{11.5} - 10^{12.0}$ cm$^{-3}$, illuminated by the ionizing radiation corresponding to the ionization parameter of $U = 10^{-3.0} - 10^{-2.5}$ \citep[][the latter is referred to as Paper I hereafter]{matsuoka05, matsuoka07}. Now that the general properties of the \ion{O}{1} and \ion{Ca}{2} emitting gas, as a whole, are thus being revealed, we should turn our eyes to the following question: how are these gas properties related to the quasar characteristics, such as redshift and luminosity? Here we present results of a compilation of the emission-line fluxes of \ion{O}{1} $\lambda$8446, \ion{O}{1} $\lambda$11287, and the near-IR \ion{Ca}{2} triplet ($\lambda$8498, $\lambda$8542, and $\lambda$8662) observed in 11 quasars. The quasars are distributed over wide ranges of redshift (0.06 $\le z \le$ 1.08) and of luminosity ($-29.8 \le M_{B} \le -22.1$), thus providing a useful sample to track the line-emitting gas properties in various quasar environments. The observational and theoretical data used in this work are described in \S \ref{sec:data}, and results and discussion appear in \S \ref{sec:results}. Our conclusions are summarized in \S \ref{sec:summary}. | } \subsection{A Picture of the Dust-Free Emission Region} We plot the observed values of $n$(\ion{Ca}{2})/$n$(\ion{O}{1} $\lambda$8446), \ion{O}{1} $n$($\lambda$11287)/$n$($\lambda$8446), and EW (\ion{O}{1} $\lambda$8446) as functions of the redshift and of the $B$-band absolute magnitude $M_{B}$ in Figure \ref{vsprop2}. One of the remarkable results is found in the top panels; the \ion{Ca}{2}/\ion{O}{1} $\lambda$8446 ratio is hardly dependent on redshift or luminosity over the plotted range, while the ratio is predicted to be very sensitive to the density of the line-emitting gas in the photoionization models. We show the model predictions on the $n$(\ion{Ca}{2})/$n$(\ion{O}{1} $\lambda$8446) -- \ion{O}{1} $n$($\lambda$11287)/$n$($\lambda$8446) plane in Figure \ref{oioi} ({\it left}), as well as the observed values (see also Fig. 7 in Paper I\footnote{ Note that Figure 7 in Paper I shows the predictions of Model 1, which is not adopted in this paper since the assumed microturbulent velocity, $v_{\rm turb}$ = 0 km s$^{-1}$, could not reproduce the observed \ion{Fe}{2} UV emissions (see \S \ref{data_model}). However, Model 3 adopted in this work predict very similar results to those shown in Figure 7 regarding the \ion{O}{1} and \ion{Ca}{2} emissions, while the whole pattern of contour is slightly ($\sim$ 0.5 dex) shifted to the high-density regime; the best-fit parameters in Model 1 are ($n_{\rm H}$, $U$) = (10$^{11.5}$ cm$^{-3}$, 10$^{-3.0}$). }). The gas density $n_{\rm H}$ and the ionization parameter $U$ are changed around the reference grid point, ($n_{\rm H}$, $U$) = (10$^{12.0}$ cm$^{-3}$, 10$^{-2.5}$), over 2 orders of magnitude in both parameters. It is clearly seen that the predicted $n$(\ion{Ca}{2})/$n$(\ion{O}{1} $\lambda$8446) ratio increases monotonically with the increased gas density, and that all the observed values are marked with the density in the vicinity of the reference point, log $n_{\rm H}$ = 12.0. Thus the similar density of the line-emitting gases are strongly indicated for the quasars distributed over these redshift and luminosity ranges. The values of EW (\ion{O}{1} $\lambda$8446), both observed and calculated with the density of log $n_{\rm H}$ = 12.0, are shown in Figure \ref{oioi} ({\it right}) as a function of the \ion{O}{1} $n$($\lambda$11287)/ $n$($\lambda$8446) ratio. It shows that the predictions of EWs are also consistent with the observed data when log $n_{\rm H}$ = 12.0 and a covering fraction (cf) of the line-emitting gas as seen from the central energy source of 0.2 -- 0.5 are assumed. Note that the covering fraction could be much smaller if we assumed oxygen overabundance relative to the solar value. On the other hand, a scatter of the observed data in the \ion{O}{1} $n$($\lambda$11287)/ $n$($\lambda$8446) axis seems to be related to the diversity of the ionization parameter (Fig. \ref{oioi}, {\it left}). Note that the diversity of other parameters, such as microturbulent velocity, gas column density, and chemical composition, could not explain this diagram since they significantly alter the $n$(\ion{Ca}{2})/$n$(\ion{O}{1} $\lambda$8446) ratio, rather than \ion{O}{1} $n$($\lambda$11287)/ $n$($\lambda$8446), and thus they are unable to explain the observed similarity of the former ratios (Paper I). It is also quite unlikely that the diversity of these parameter values is balanced out by the fine-tuned density in such a way that the $n$(\ion{Ca}{2})/$n$(\ion{O}{1} $\lambda$8446) ratio is always kept to be $\sim$1.0, unless these lines are the dominant heating or cooling sources of the emission region. As with the $n$(\ion{Ca}{2})/$n$(\ion{O}{1} $\lambda$8446) ratio, \ion{O}{1} $n$($\lambda$11287)/ $n$($\lambda$8446) is not clearly dependent on the redshift or luminosity (Fig. \ref{vsprop2}, {\it middle left and middle right}). The above arguments invoke a picture of the line-emitting gases in quasars that have similar densities and are located at regions exposed to various ionizing radiation fluxes. It would be a consequence of the difference in distance to the central continuum source and/or in the intrinsic luminosity of the quasars. Note that it is in clear contrast to the well-studied case of H$\beta$, whose emission regions in AGNs are known to be characterized by similar ionization parameters. In fact, reverberation mapping results for H$\beta$ show the emission region size ($r$) -- luminosity ($L$) relation of $r \propto L^{0.5}$, which is consistent with the constant ionization parameter regime \citep{peterson02, bentz06}. Such a situation has long been expected in order to account for the remarkably similar AGN spectra over a broad range of luminosity, and was incorporated into the locally optimally emitting cloud (LOC) model suggested by \citet{baldwin95}, that is, that the BELR is composed of gas with widely distributed physical parameters and each emission line arises from its preferable environment. On the other hand, the case in \ion{O}{1} and \ion{Ca}{2} lines apparently indicates that the location of the emission region is not radiation-selected. In line with the above arguments, we found a clear difference of the velocity-luminosity relation between \ion{O}{1} and H$\beta$; the measured \ion{O}{1} line widths are plotted versus $M_B$ in Figure \ref{mag_width}, which shows that the \ion{O}{1} line widths are remarkably similar, concentrated around 1500 -- 2000 km s$^{-1}$, over more than 3 orders of magnitude in the $B$-band luminosity. On the other hand, those for the hydrogen Balmer lines usually have a large scatter as shown by, e.g., \citet{kaspi00}; their sample of 34 AGNs, spanning over 4 orders of magnitude in continuum luminosity, has the line widths of 1000 -- 10,000 km s$^{-1}$. Such a trend is also indicated by \citet{persson88}, who reported that while the correlation between FWHM (\ion{O}{1}) and FWHM (H$\beta$) is good for the small FWHM regime, the \ion{O}{1} lines grow systematically narrower than H$\beta$ at large line width. \citet{rodriguez02a} conducted a detailed study of the near-IR emission line profiles in NLS1s, and found that \ion{O}{1}, \ion{Ca}{2}, and \ion{Fe}{2} lines are systematically narrower than the broad components of other low-ionization lines such as hydrogen Paschen lines and \ion{He}{1} $\lambda$10830. They also argued that these lines are produced in the outermost portion of the BELR, since their widths are just slightly broader than those of [\ion{S}{3}] $\lambda$9531 which they assumed is formed in the inner portion of the narrow emission line region (NELR). While the scattered widths of \ion{H}{1} lines could be interpreted as a consequence of the radiation-selected locations of the emitting gases, regardless of the gas kinematics, the remarkable similarity of the \ion{O}{1} line widths might imply the kinematically determined emission regions. In such a situation, the diversity of the ionization parameters, as discussed above, would be an inevitable result. As stated by \citet{persson88}, there is clearly interesting information which could be deduced from studies of the \ion{O}{1} and \ion{H}{1} line profiles; especially, reverberation mapping of these \ion{O}{1} and \ion{Ca}{2} lines would be a powerful tool to reveal the underlying physics. \subsection{A Picture of the Dusty Emission Region} The widely-accepted theory of the BELR describes its outer edge, where the \ion{O}{1} and \ion{Ca}{2} emission lines are likely to be formed, set by the dust sublimation \citep[e.g.,][]{laor93, netzer93}. If we accept this picture, the dust grains are possibly mixed in the line-emitting gas and suppress the \ion{Ca}{2} emission through the substantial Ca depletion. % Such a situation is in fact reported for the NELR by the absence or significant weakness of the observed [\ion{Ca}{2}] $\lambda$7291 line \citep{kingdon95, villar97}. \citet{ferguson97} presented the LOC model calculations of the narrow emission lines and argued that Ca is depleted relative to the solar value by factors of 3 -- 160. It is hard to see an evidence of the dust presence in the emission region from our results, since Figure \ref{oioi} appears to show that our dust-free models successfully reproduce the observations. However, it is noteworthy that % it is quite difficult to account for % the observed data at \ion{O}{1} $n$($\lambda$ 11287)/$n$($\lambda$ 8446) $<$ 0.4 by the models with log $n_{\rm H}$ = 12.0. The problem is that such small values of \ion{O}{1} $n$($\lambda$ 11287)/$n$($\lambda$ 8446) could only be reproduced with the higher densities than log $n_{\rm H}$ = 12.0 so that \ion{O}{1} $\lambda$8446 emission is exclusively enhanced by the collisional excitation, while such a dense gas produces intense \ion{Ca}{2} emission that is much stronger than observed. One can clearly see this trend in Figure \ref{oioi} ({\it left}). If we assumed significant Ca depletion in the line-emitting gas, these difficulties are naturally resolved since it significantly suppresses the otherwise intense \ion{Ca}{2} emissions. For example, the data point representing (log $n_{\rm H}$, log $U$) = (13.0, $-$2.5) in Figure \ref{oioi} ({\it left}) could provide a plausible model for the observations at \ion{O}{1} $n$($\lambda$ 11287)/$n$($\lambda$ 8446) $<$ 0.4 if the \ion{Ca}{2} emission is suppressed by a factor of a few times 10.\footnote{ The predicted EW for the model with (log $n_{\rm H}$, log $U$) = (13.0, $-$2.5) is EW (\ion{O}{1} $\lambda$8446) = 30 \AA, which is consistent with the observations if rather large covering factors of 0.5 -- 1.0 and/or oxygen overabundance relative to the solar value are assumed. } The dust grains present in the line-emitting gas might also give the natural explanation to the observed lack of some UV emission lines relative to their optical or near-IR counterparts. \citet{ferland89} mentioned the possibility of the dust survival in the BELR gas in order to explain the extreme weakness of the observed \ion{Ca}{2} $\lambda$3934 and $\lambda$3639 lines relative to the near-IR triplet. At least a part of the long-standing \ion{Fe}{2} UV/opt problem, in which the observed ratios of \ion{Fe}{2} UV flux to the optical flux fall far below the photoionization model predictions \citep[see, e.g.,][]{baldwin04}, could also be explained. However, it should be noted that the dust would affect the line formation processes in a very complicated manner, which should be precisely addressed when discussing the specific lines; for the \ion{O}{1} and \ion{Ca}{2} lines, one of the most apparent effects as well as the Ca depletion would be the destruction of Ly$\beta$ photons which otherwise excite \ion{O}{1} atoms, thus the suppression of the \ion{O}{1} emissions. The resultant line flux ratios could be much different from those derived from the simple speculations. } We have compiled the emission-line fluxes of \ion{O}{1} $\lambda$8446, \ion{O}{1} $\lambda$11287, and the near-IR \ion{Ca}{2} triplet ($\lambda$8579) observed in 11 quasars. The quasars are distributed over wide ranges of redshift (0.06 $\le z \le$ 1.08) and of luminosity ($-29.8 \le M_B \le -22.1$), thus providing a useful sample to track the line-emitting gas properties in various quasar environments. The measured line strengths and velocities, as functions of the quasar properties, were analyzed with photoionization model calculations. Our findings and conclusions are as follows: 1. There is no sign of a significant change in the flux ratios of the \ion{Ca}{2} triplet and \ion{O}{1} $\lambda$8446 over the redshift and luminosity ranges studied here. It strongly indicates similar gas densities in the line-emission region from quasar to quasar. 2. The observed scatter of the \ion{O}{1} $\lambda$11287/$\lambda$8446 ratios appears to be related to the diversity of the ionization parameter, while the ratio is not clearly dependent on the redshift or luminosity. Combined with the similarity of the \ion{Ca}{2}/\ion{O}{1} $\lambda$8446 ratios, it invokes the picture of the line-emitting gases in quasars that have similar densities and are located at regions exposed to various ionizing radiation fluxes. 3. The \ion{O}{1} line widths are remarkably similar from quasar to quasar over more than 3 orders of magnitude in luminosity. It might imply a kinematically determined location of the line emission region and is in clear contrast to the case of \ion{H}{1} lines, whose emission region is considered to be radiation-selected. 4. If we accept that the \ion{O}{1} and \ion{Ca}{2} emission lines are formed at the outer edge of the BELR and that the outer edge is set by the dust sublimation, the line-emitting gas is possibly mixed with the dust grains. In fact such a situation may better reproduce the observations than the dust-free case through the significant Ca depletion. | 7 | 10 | 0710.2954 |
0710 | 0710.5603_arXiv.txt | We review the current status of accelerator, direct and indirect Dark Matter (DM) searches, focusing on the complementarity of different techniques and on the prospects for discovery. After taking a census of present and upcoming DM-related experiments, we review the motivations to go beyond an "accelerator-only" approach, and highlight the benefits of multidisciplinarity in the quest for DM. | The evidence for non-baryonic dark matter is compelling at all observed astrophysical scales~\cite{Bergstrom:2000pn,Bertone:2004pz}. Although alternative explanations in terms of modified gravity (see Ref.~\cite{Bekenstein:2004ne} for a relativistic theory of the MOND paradigm) cannot be ruled out, they can hardly be reconciled with the most recent astrophysical observations~\cite{Clowe:2006eq} without requiring additional matter beyond the observed baryons (e.g. Ref.~\cite{Feix:2007zm} and references therein). It is therefore natural to ask {\it how can we identify the nature of DM particles?}. We review here the main strategies that have been devised to attack this problem, namely accelerator, direct and indirect searches, focusing on the interplay between them and on their complementarity. In fact, a tremendous theoretical and experimental effort is in progress to clarify the nature of DM, mostly devoted, but not limited, to searches for Weakly Interacting Massive Particles (WIMPs), that achieve the appropriate relic density by {\it freezing-out} of thermal equilibrium when their self-annihilation rate becomes smaller than the expansion rate of the Universe. The characteristic mass of these particles is $\cal{O}$$(100)$ GeV, and the most representative and commonly discussed candidates in this class of models are the supersymmetric neutralino, and the B$^{(1)}$ particle, first excitation of the hypercharge gauge boson, in theories with Universal Extra Dimensions. A tentative census of present and upcoming DM experiments (WIMPs only) is shown in fig. 1. Shown in the figure are: two particle accelerators, viz. the Tevatron at Fermilab, and the upcoming Large Hadron Collider (LHC) at CERN; the many direct detection experiments currently taking data or planned for the near future, along with the names of the underground laboratories hosting them; high-energy neutrino telescopes; gamma-ray observatories; gamma-ray and anti-matter satellites. Light blue points denote gamma-ray experiments that are not directly related to indirect DM searches, as DM signals would be typically produced at energies below their energy threshold. Nevertheless, they may turn out to be useful to discriminate the nature of future unidentified high-energy gamma-ray sources. Three satellites are shown in the inset of figure 1: PAMELA, an anti-matter satellite that has already been launched and is expected to release the first scientific data very soon. ; GLAST, a gamma-ray satellite that is scheduled for launch in early 2008; and AMS-02, anti-matter satellite that should be launched in the near future. We will discuss below the prospects for detecting DM with the various experiments shown in fig. 1, and we will focus our attention on the complementarity of the various detection strategies. The paper is organized as follows: we first discuss accelerator searches, and show that although the LHC has the potential to make discoveries of paramount importance for our understanding of DM, it may not be able to solve all problems. In Section 3 we discuss the information that can be extracted from direct detection experiments, in case of positive detection. Section 4 is then dedicated to indirect searches, and to the question of what astrophysical observations can tell us about the nature of DM, and how to combine this information with all other searches. \begin{figure*} \centering \includegraphics[width=\textwidth]{dmmap2.eps} \caption{2007 census of present and upcoming Dark Matter-related experiments. Black points denote the location of high energy neutrino telescopes; Dark-blue points are for gamma-ray Air Cherenkov Telescopes, while light-blue points are for other ground-based gamma-ray observatories. Red points are for underground laboratories hosting existing and upcoming direct detection experiments. Yellow points show the location of the Fermilab's Tevatron, and the upcoming Large Hadron Collider at CERN.} \label{DMmap} \end{figure*} | 7 | 10 | 0710.5603 |
|
0710 | 0710.2215_arXiv.txt | {Sequences of Doppler images of the young, rapidly rotating late-type stars AB Dor and LQ Hya show that their equatorial angular velocity and the amplitude of their surface differential rotation vary versus time. Such variations can be modelled to obtain information on the intensity of the azimuthal magnetic stresses within stellar convection zones. We introduce a simple model in the framework of the mean-field theory and discuss briefly the results of its application to those solar-like stars. } | \noindent Doppler imaging techniques allow us to measure the surface differential rotation in rapidly rotating late-type stars by tracking the longitudinal motion of starspots located at different latitudes (Collier Cameron 2007). Spe\-ci\-fi\-cal\-ly, the surface angular velocity $\Omega$ at colatitude $\theta$ is assumed to be given by a solar-like law: \begin{equation} \Omega(\theta) = \Omega_{\rm eq} - d\Omega \cos^{2} \theta, \end{equation} where $\Omega_{\rm eq}$ is the equatorial angular velocity and $d\Omega$ is the pole-equator angular velocity difference. $\Omega_{\rm eq}$ and $d\Omega$ can be measured by fitting the shear of starspots along sequences of photospheric images covering successive rotations. Alternatively, $\Omega$ and $d\Omega$ can be included as free parameters in a code that reproduces the line profile distortions due to starspots, thus obtaining their values by a suitable $\chi^{2}$ minimization when a sufficiently long time series of line profiles is available. The long-term monitoring of surface differential rotation of the two late-type dwarfs AB Doradus and LQ~Hydrae shows that their equatorial angular velocity and surface shear are functions of the time (see Donati et al. 2003a; Jeffers et al. 2007). It is interesting to note that the variations of $\Omega_{\rm eq}$ and $d\Omega$ are compatible with an internal angular velocity uniform on cylindrical surfaces co-axial with the rotation axis (Donati et al. 2003a). | We modelled the observed time variation of the differential rotation in AB~Dor and LQ~Hya under the hypotheses that the azimuthal Maxwell stresses rule the changes of their surface shear and the internal angular velocity depends only on the distance from the rotation axis (Taylor-Proudman regime). We obtained that the average intensity of the mean field Maxwell stres\-ses is $ |B_{s} B_{\varphi} | \sim 0.03 - 0.14 $ T$^{2}$, implying azimuthal mean fields $B_{\varphi} \sim (3-10)$ T for $B_{\rm s} \sim 0.01$ T. Similar Maxwell stresses are obtained if the magnetic torque is assumed to be confined within the overshoot layer $0.67 \leq (r/R) \leq 0.71$ and no restrictions are imposed on the internal rotation law. It is interesting to note that azimuthal magnetic fields of $3-10$ T, occupying a sizeable fraction of the convection zone, have been invoked to explain orbital period changes observed in late-type active binaries (Lanza, Rodon\`o \linebreak \& Rosner 1998; Lanza \& Rodon\`o 2004; Lanza 2005, \linebreak 2006b). An $\alpha$-effect related to an instability of the magnetic field itself (e.g., magnetic buoyancy instability, Brandenburg \& Schmitt 1998; or magneto-rotational instability, R\"udiger et al. 2007) seems to be necessary to produce such super-equi\-par\-ti\-tion fields, possibly acting in combination with differential rotation. The energy dissipated by turbulence, estimated according to standard mixing-length arguments, may exceed stellar luminosity in the case of the largest surface shear observed in LQ Hya. However, the thermal equilibrium of the convection zone can be significantly affected only if those large shear episodes last more than $\sim 10-20$ per cent of the time. Note also that a mixing-length estimate for the turbulent viscosity may not be appropriate for a rapidly rotating star (see Lanza 2006a for details). In the present work, we adopted a point of view analogous to that of Covas et al. (2005) who modelled the variation of stellar differential rotation considering only the \linebreak torque exerted by the Lorentz force and neglecting the roles of meridional flow and $\Lambda$-effect. As a matter of fact, it is difficult to evaluate the perturbations of the meridional flow and of the Reynolds stresses produced by the magnetic field because they depend critically on the approximations made in the treatment of stellar turbulence in a rotating star. Nevertheless, alternative models for the variation of differential rotation have been investigated, such as those based on a time-dependent component of the meridional flow (Rempel 2006, 2007) or the magnetic quenching of the $\Lambda$-effect (R\"udiger et al. 1986). The present approach can be generalized to obtain amplitudes of the perturbations of the corresponding terms in Eq.~(\ref{tau}), but we shall not pursue this application here (see Lanza 2007). Finally, it is interesting to note that some inference on the amplitude of variation of the surface shear in the case of very active stars can be obtained not only by means of Doppler imaging techniques, but also by an appropriate ana\-ly\-sis of their long-term wide-band photometry (e.g., \linebreak Rodon\`o et al. 2001; Messina \& Guinan 2003). For example, in the case of LQ~Hya, it is worth comparing the Doppler imaging results by Donati et al. (2003a) with the photometrically determined seasonal rotation periods by Kov\'ari et al. (2004). | 7 | 10 | 0710.2215 |
0710 | 0710.0210_arXiv.txt | The unified dark energy and dark matter model within the framework of a model of a continuous medium with bulk viscosity (dark fluid) is considered. It is supposed that the bulk viscosity coefficient is an arbitrary function of the Hubble parameter. The choice of this function is carried out under the requirement to satisfy the observational data from recombination ($z\approx 1000$) till present time. | Modelling of an accelerated expansion of the present Universe lies on the way of creation of phenomenological models which may explain the observational data on one set of parameters and compare them with predictions of the models on other set. For example, a theoretical model adapts for the correct description of the acceleration of the Universe in an accessible interval $z$. Further, the results of modelling are extrapolated for large $z$ which are not accessible for observations yet. The corresponding cosmological scenario defines growth of the large-scale structure which determines the present-day fluctuations of the microwave background radiation. Certainly, the models under consideration should not contradict the available observational data within the framework of general relativity in a field of its applicability. For the specified purposes a number of cosmological models successfully applied in the past in the theory of the early Universe is used. These are cosmological models with various scalar and non-scalar fields filling the space together with cold dark matter (see, e.g., the reviews \cite{Sahni}). Cosmological models of the present accelerated Universe within the framework of high-order theories of gravity (HOTG) are also quite popular \cite{Carrol}. A number of models have recently been suggested \cite{Ren, Brevik,Odin} which describe the present Universe with use of models of a continuous medium in the presence of bulk viscosity. Consideration of effects of viscosity within the framework of HOTG was also carried out \cite{Brevik1}. Note that such models were well-known in the theory of the early Universe (see, for example, \cite{Murphy,Barrow}). In particular, in Ref. \cite{Barrow} a few exact solutions with the constant bulk viscosity coefficient and with the bulk viscosity being an arbitrary power function of energy density were obtained. In Ref. \cite{Ren} the model of viscous dark fluid is considered. The main result of this paper is the model with the constant bulk viscosity coefficient. The model fits the observational data on luminosity at an acceptable level. In Ref. \cite{Brevik} the models both with the constant bulk viscosity coefficient and the bulk viscosity linearly proportional to the Hubble parameter are examined. The question about influence of viscosity on presence of a singularity in the model in the future (the so-called Big Rip) is investigated. In this paper we consider a model of "viscous dark fluid" with the bulk viscosity coefficient $\mu(H)$ which depends on the Hubble parameter arbitrarily. Unlike Ref. \cite{Ren}, comparison of the model with the observational data is not restricted to the observational data on luminosity. The model is being compared with results of observations on change of the deceleration parameter $q$ and values of the Hubble parameter in the range $2>z>0$. It will be shown below that the model with the constant bulk viscosity coefficient does not provide a good description for $q(z)$ and $H(z)$ which follow from the observations. We propose such a dependence $\mu(H)$ which is adequate to the mentioned observations. The proposed model is extrapolated for $z$ beyond the specified range $2>z>0$. | The bulk viscosity in our model is an example of a dynamic $\Lambda$-term. However, our model is close to the $\Lambda$CDM model. As it was rightly noted in \cite{Ren}, the model with the viscosity does not give possibility to divide the true dust filling the Universe, and dark matter generated by the bulk viscosity. That is why it is difficult to introduce a phenomenological equation of state $p=w \varepsilon$ which is often used for interpretation of the observational data. Influence of the bulk viscosity on formation of the large-scale structure of the Universe demands special examination. But taking into account that the viscosity is being ''involved`` at rather small $z\approx 2$, it is possible to expect its influence on the dynamics of galactic clusters only at later non-linear stage. | 7 | 10 | 0710.0210 |
0710 | 0710.5435_arXiv.txt | {\object{NRAO~150} --a compact and bright radio to mm source showing core/jet structure-- has been recently identified as a quasar at redshift $z=1.52$ through a near-IR spectral observation.} {To study the jet kinematics on the smallest accessible scales and to compute the first estimates of its basic physical properties,} {we have analysed the ultra-high-resolution images from a new monitoring program at 86~GHz and 43~GHz with the GMVA and the VLBA, respectively. An additional archival and calibration VLBA data set, covering from 1997 to 2007, has been used.} {Our data shows an extreme projected counter-clock-wise jet position angle swing at an angular rate of up to $\approx 11^{\circ}/\rm{yr}$ within the inner $\approx 31$~pc of the jet, which is associated with a non-ballistic superluminal motion of the jet within this region.} {The results suggest that the magnetic field could play an important role in the dynamics of the jet in \object{NRAO~150}, which is supported by the large values of the magnetic field strength obtained from our first estimates. The extreme characteristics of the jet swing make \object{NRAO~150} a prime source to study the jet wobbling phenomenon.} | \label{int} An increasing number of jets in active galactic nuclei (AGN) have been reported to show either regular or irregular swings of the innermost jet structural position angle in the plane of the sky (e.g., in \object{OJ~287}, Tateyama \& Kingham~\cite{Tat04}; in \object{3C~273}, Savolainen et al.~\cite{Sav06}; in \object{3C~345}, Lobanov \& Roland~\cite{Lob05}; in \object{BL~Lac}, Stirling et al.~\cite{Sti03}; in \object{S5~0716+71}, Bach et al.~\cite{Bac05}). Time scales between 2 and 15 years and structural position-angle oscillations with amplitudes from $\sim 25^{\circ}$ to $\sim 45^{\circ}$ are typical for the reported cases. We will call this phenomenon \emph{jet wobbling} hereafter. Parsec scale AGN jet curvatures and helical-like structures also at larger distances from the central engine are also believed to be triggered by changes in direction at the jet ejection nozzle (e.g., in \object{3C~84}, Dhawan et al.~\cite{Dha98}). The physical origin for the observed jet wobbling is still poorly understood. Among the various possibilities, regular precession of the accretion disk is frequently used for modeling at present. Most AGN precession models are driven either by a companion super-massive black hole or another massive object inducing torques in the accretion disk of the primary (e.g., Lister et al.~\cite{Lis03}, for \object{4C~+12.50}; Stirling et al.~\cite{Sti03}, for \object{BL~Lac}; Caproni \& Abraham~\cite{CapAb04} for \object{3C~120}) or by the Bardeen-Peterson effect (e.g., Liu \& Melia~\cite{Liu02}; Caproni et al.~\cite{Cap04}). However, other AGN scenarios that have yet to be explored extensively, such as the orbital motion of the jet nozzles (also involving binary systems) or other kinds of more erratic disk/jet instabilities (e.g., similar to those thought to produce the quasi periodic oscillations [QPO] in X-ray binaries), can not be ruled out yet. Note that, in support of these erratic instabilities, it is still under debate whether the observed jet wobbling is strictly periodic or not (see Mutel \& Denn~\cite{Mut05} for the case of \object{BL~Lac}). There is still no paradigm to explain the phenomenon of jet wobbling in AGN, but it is rather likely that, as it is triggered in the innermost regions of the jets, it must be tied to fundamental properties of the inner regions of the accretion system. Hence, there is ample motivation to study the jet wobbling phenomenon to place our understanding of the jet triggering region and the super--massive accretion systems on firmer ground. VLBI observations at millimetre wavelengths are a powerful technique to image the innermost regions of AGN jets -which are self-absorbed at longer wavelengths- with the highest angular resolutions; $\sim 50\,\mu$as at 86~GHz (3.5~mm) and $\sim 0.15$\,mas at 43~GHz (7~mm). Here, we report the discovery of an extreme case of jet swing in the quasar \object{NRAO~150}, through the first ultra-high-resolution VLBI set of images obtained from this source at 86~GHz and 43~GHz. \object{NRAO~150} is a strong radio-mm source. At radio wavelengths, on VLBI scales, \object{NRAO~150} displays a compact core plus a one-sided jet extending up to $r \apgt 80$~mas with a jet structural position angle (PA) of $\sim 30^\circ$ (e.g., Fey \& Charlot~\cite{Fey00}). \object{NRAO~150} was not detected by the early optical surveys, most probably due to obscuration through the Milky Way (Galactic latitude $= -1.6^\circ$). Almost 40~yr after its discovery at radio wavelengths (Pauliny-Toth, Wade \& Heeschen~\cite{Pau66}), \object{NRAO~150} has been identified as a quasar at redshift $z=1.52$ through a near-IR spectroscopic project (Acosta-Pulido et al. in prep.). Throughout this paper we assume $H_{\circ} = 72$~km~s$^{-1}$~Mpc$^{-1}$, $\Omega_{m}=0.3$, and $\Omega_{\Lambda}=0.7$. Under these assumptions, the luminosity distance of \object{NRAO~150} is $d _{L}=11025$~Mpc, 1~mas corresponds to $8.5$~pc in the frame of the source, and an angular proper motion of 1~mas/yr translates into a speed of $69.4~c$. These parameters are used here for the first time to make quantitative estimates of the basic physical properties of the jet in \object{NRAO~150}. | \label{concl} We have reported the results from the first 86~GHz and 43~GHz VLBI monitoring program of the recently identified quasar \object{NRAO~150}. Our observations reveal {\it a)} a large projected misalignment of the jet by $>100^{\circ}$ within the inner 0.5~mas to 1~mas from the core, {\it b)} an extremely fast counter-clockwise rotation of the projected jet axis at a rate of $\sim6^{\circ}/\rm{yr}$ to $\sim11^{\circ}/\rm{yr}$, {\it c)} non-ballistic superluminal motions with mean speeds from 2.3~$c$ to 3.3~$c$ within the inner 0.5~mas from the core (deprojected distance $\approx 31$~pc), {\it d)} transverse (non-radial) speeds of 2.7~$c$, 1.4~$c$, and 2.0~$c$ for Q1, Q2 and Q3, respectively, {\it e)} an extreme case of \emph{superluminal non-ballistic jet swing}, and {\it f)} the first approximations to quantitative estimates of the basic physical properties of the jet in \object{NRAO~150}: $\delta \approx 6$, $B\approx 0.7$\,G, $\gamma \approx 4$, and $\phi \approx 8^{\circ}$. Whereas the ultimate origin of the jet swing must be an intrinsic change of the direction of the inner jet axis (either caused by changes produced at the injection region or by interaction with the medium surrounding the jet), possible causes for the non-ballistic nature of the emitting flow are either a small inertia of the jet compared to that of the impacted ambient medium, or by the fact that we are observing a jet instability propagating downstream. However, the superluminal non-ballistic motion of the jet features in \object{NRAO~150} is perhaps too fast and systematic along tens of parsecs (deprojected, see above) to be induced solely by the ambient medium. In addition, the perturbation that must be produced either by the impact with such medium or by the changes at the injection region should imply the growth of disruptive instabilities (e.g., Perucho et al.~\cite{Per06}), in contrast to the remarkable collimation of the jet up to the kiloparsec scale (see Fig.~1). Mizuno et al.~(\cite{Miz07}) have shown that a magnetic field with the appropriate configuration helps keeping a jet collimated against the growth of instabilities. Hence, this possibility, together with the high value for the magnetic field intensity estimated in Section~\ref{phys-par}, suggests that the magnetic field could play an important role in the dynamics of the jet in \object{NRAO~150}. It is still unclear whether the reported change of the direction of ejection in \object{NRAO~150} is related to a regular (strictly periodic or not) behaviour or to a single event. This, together with the still unknown nature of the moving jet features (either propagating curvatures or shocks in a bent jet), does not allow us to specify the ultimate origin of this phenomenon in the source. This will be the matter of future studies. Nevertheless, the extreme characteristics of the jet swing make \object{NRAO~150} a prime source for future studies of the origin of jet wobbling. | 7 | 10 | 0710.5435 |
0710 | 0710.0817_arXiv.txt | The fast rotating star CU Virginis is a magnetic chemically peculiar star with an oblique dipolar magnetic field. The continuum radio emission has been interpreted as gyrosyncrotron emission arising from a thin magnetospheric layer. Previous radio observations at 1.4~GHz showed that a 100\% circular polarized and highly directive emission component overlaps to the continuum emission two times per rotation, when the magnetic axis lies in the plane of the sky. This sort of radio lighthouse has been proposed to be due to cyclotron maser emission generated above the magnetic pole and propagating perpendicularly to the magnetic axis. Observations carried out with the Australia Telescope Compact Array at 1.4 and 2.5~GHz one year after this discovery show that this radio emission is still present, meaning that the phenomenon responsible for this process is steady on a timescale of years. The emitted radiation spans at least 1 GHz, being observed from 1.4 to 2.5~GHz. On the light of recent results on the physics of the magnetosphere of this star, the possibility of plasma radiation is ruled out. The characteristics of this radio lighthouse provides us a good marker of the rotation period, since the peaks are visible at particular rotational phases. After one year, they show a delay of about 15 minutes. This is interpreted as a new abrupt spinning down of the star. Among several possibilities, a quick emptying of the equatorial magnetic belt after reaching the maximum density can account for the magnitude of the breaking. The study of the coherent emission in stars like CU~Vir, as well as in pre main sequence stars, can give important insight into the angular momentum evolution in young stars. This is a promising field of investigation that high sensitivity radio interferometers such as SKA can exploit. | CU Virginis (=HD124224) is an A-type magnetic chemically peculiar star (MCP) with a rotational period of 0.52 days, one of the shortest for this class of stars. As observed in all MCP stars, the variability of the light curve is correlated with the spectroscopic variations \citep{deutsch, hardie} and with the effective magnetic field \citep{borra80}. All those characteristics can be explained in the framework of the oblique rotator model, where the axis of the dipolar magnetic field is tilted with respect to the rotational one \citep{babcock} and abundance of elements is not homogeneously distributed over the stellar surface. The observed variabilities are thus consequence of the stellar rotation. \begin{figure*} \includegraphics[width=17cm]{spettri.ps} \caption{Dynamical spectra of CU~Vir during the two days of observations (left panels: May 29, 1999) at 2.5~GHz (upper panels) and 1.4~GHz (lover panels). Some spectral channels have been removed, reducing the bandpass. The spectral resolution is 8~MHz; the spectra have been smoothed in time with a window 2 minutes wide. Strong enhancements of the radio emission are evident at two phases at 1.4~GHz, while only one peak is visible at 2.5~GHz. } \label{spettri} \end{figure*} MCP stars are in general slow rotators in comparison with normal B and A-type main sequence stars. This behaviour can be explained as the result of the action of a magnetic breaking. But, at the present, only two MCP stars have been found to increase their rotational period; they are 56 Ari and CU Vir. While the former shows a continuous breaking down at the rate of few seconds per century \citep{adelman}, CU Vir has been subject of an abrupt increase of the rotational period \citep{pyper}. From the analysis of photometric light curves collected over 40 years, it seems that a change of period of about two seconds occurred abruptly in 1984. The mechanism responsible for this event is still under debate, and precise timing of the rotation is needed in order to detect any further slowing down of the star. Radio emission has been observed in CU Vir \citep{leo94}. The radio spectrum is quite flat and extends up to mm wavelengths \citep{leo2004} with an high degree of circular polarization \citep{leo96}. The variabilities of total intensity and polarization are both correlated with the rotation of the star, suggesting we are in presence of gyrosyncrotron emission from an optically thick source. The anisotropic stellar wind inferred from spectral line variations in MCP star \citep{shore}, associated with the magnetic field justifies the radio emission. \citet{trig04} developed a three-dimensional model to explain the radio emission from MCP stars. The dipolar magnetic field interacts with the stellar wind, that can freely escape from the polar regions outside the Alfv\'en surface (outer magnetosphere) but remains trapped in the equatorial belt (inner magnetosphere). Ionized particles flowing out in the transition region between outer and inner magnetosphere, called middle magnetosphere, stretch the magnetic field lines just outside the Alfv\'en radius and open the field generating current sheets, where particles are accelerated up to relativistic energies. They eventually propagate back toward the stellar magnetic poles following the field lines and radiating for gyrosyncrotron process. As the magnetic field intensity increases traveling to the star, they are reflected back outward by the magnetic mirroring effect and are definitively lost from the magnetosphere. This model, used to explain the observed fluxes and variability of MCP stars (HD37479 and HD37017) has been successfully applied to CU Vir by \citet{leto06}. On the basis of multiwavelengths radio observations, important physical parameters of the stellar magnetosphere, as the Alfv\'en radius ($12-17\,R_\ast$) and the mass loss (about $10^{-12}M_\odot \rmn{yr}^{-1}$) have been inferred. A further observational evidence supporting the picture outlined above is the discovery of coherent, highly directive, 100\% polarized radio emission at 1.4~GHz \citep{trig00}. The two peaks of radio emission have been observed at the rotational phases when the magnetic axis of the dipole is perpendicular to the line of sight. The two peaks have been observed in three observing runs spanning 10 days, indicating that the emission mechanism is persistent at least in timescales of weeks. This has been interpreted as Electron Cyclotron Maser Emission from a population of electrons accelerated in the current sheets at the Alfv\'en point, that developed a loss cone anisotropy after the mirroring and masing in a direction almost perpendicular to the motion, so to the field lines, just above the magnetic pole. In this paper we present radio observations of CU Vir at 1.4 and 2.5~GHz carried out with the aim to confirm the maser emission and to study its spectrum. The directivity of the radiation is used to check the rotational period as the beam point toward the Earth two times per rotation. | The observations of CU Vir carried out with the ATCA at 1.4 and 2.5 GHz confirm the presence of the coherent emission already reported by \citet{trig00} after one year from the discovery, indicating that this is a steady phenomenon. The emitted radiation is visible only at rotational phases corresponding to the instant when the oblique axis of the magnetic dipole lies on the plane of the sky. This indicates the high directivity of this component of the radio emission. All those characteristics, and the fact that it is 100\% right hand circularly polarized, are in agreement with the process of electron cyclotron maser from electrons accelerated in the current sheets out of the Alfv\'en radius, propagating back to the stellar polar caps and reflected outward by the converging magnetic field. The lack of reflected electrons at small pitch angle, that fall in the stellar atmosphere, is the cause of the loss cone anisotropy that, in turn, generates this auroral radiation above the magnetic pole of the star. The polarization properties, the fact that the radiation is only right hand polarized, means that this process in efficient only above the north magnetic pole. This is possible as the magnetic field is not purely dipolar. The possibility of plasma radiation is ruled out since the high magnetic field strength in the region where the radiation is generated. Since the magnetosphere rotates obliquely around the rotational axis, the cyclotron maser is visible only when it points toward the Earth, like a lighthouse. In this mode it is a good marker of the rotation of the star. The analysis done shows that the peaks are delayed of about 15 minutes with respect to the observations carried out one year before, indicating a possible change of the rotational period of the star, of the order of 1 second, occurred in the period 1998--1999. A similar change of period \citep{pyper}, occurred around 1985, has been already reported. CU Vir is the unique single main sequence star with frequent abrupt spindown. Different spinning down mechanisms are discussed: the possibility of a change of the moment of inertia of the star, the continuous spindown due to the wind flowing from the Alfv\'en surface and the violent emptying of the inner magnetosphere. In this latter hypothesis the material accumulated in the closed field lines of the equatorial magnetic belt reaches a maximum density and opens the field lines in a violent event, releasing an angular momentum which may account the observed breaking. Cyclotron maser emission from stars provides important information on the magnetospheres, as it has been observed in flare stars, in dMe and brown dwarfs. In the future, when high sensitivity radio interferometers such as SKA will allow to discover more and more radio lighthouse of the same kind of CU Vir, a new possibility to study with high precision the angular momentum evolution of young main sequence and pre main sequence stars will be opened. | 7 | 10 | 0710.0817 |
0710 | 0710.1811_arXiv.txt | {}{X-ray Bright Optically Normal Galaxies (\xbongs) constitute a small but not negligible fraction of hard X-ray selected sources in recent \chandra\ and \xmm\ surveys. Even though several possibilities were proposed to explain why a relatively luminous hard X-ray source does not leave any significant signature of its presence in terms of optical emission lines, the nature of \xbongs\ is still subject of debate. We aim to a better understanding of their nature by means of a multiwavelength and morphological analysis of a small sample of these sources. } {Good-quality photometric near-infrared data (ISAAC/VLT) of four low-redshift ($z=0.1-0.3$) \xbongs, selected from the {\rm HELLAS2XMM} survey, have been used to search for the presence of the putative nucleus, applying the surface-brightness decomposition technique through the least-squares fitting program GALFIT. } {The surface brightness decomposition allows us to reveal a nuclear point-like source, likely to be responsible of the X-ray emission, in two out of the four sources. The results indicate that moderate amounts of gas and dust, covering a large solid angle (possibly 4$\pi$) at the nuclear source, combined with the low nuclear activity, may explain the lack of optical emission lines. The third \xbong\ is associated with an X-ray extended source and no nuclear excess is detected in the near infrared at the limits of our observations. The last source is associated to a close (d$\leq$ 1 arcsec) double system and the fitting procedure cannot achieve a firm conclusion.}{} | Thanks to the \chandra\ and \xmm\ surveys, the hard X-ray sky is now probed down to a flux limit where the bulk of the X-ray background is almost completely resolved into discrete sources (Hasinger et al. 2001; Alexander et al. 2003; Bauer et al. 2004; Worsley et al. 2004, 2005). Extensive programs of multiwavelength follow-up observations showed that the large majority of hard X-ray selected sources are identified with Active Galactic Nuclei (AGN) spanning a broad range of redshifts and luminosities. At variance with optically selected quasars, X-ray selected AGN are characterized by a much larger spread in their optical properties, especially for what concerns the intensity of the emission lines. Indeed, a sizable fraction of relatively luminous X-ray sources hosting an active nucleus would not have been easily recognized as such on the basis of optical observations either because associated with very faint ($R >$ 24) counterparts (e.g., Fiore et al. 2003; Mignoli et al. 2004; Civano et al. 2005) or due to the lack of AGN emission lines in their optical spectra. The latter class of sources is variously termed as ``optically-dull", ``optically normal" or \xbongs\ (X-ray Bright Optically Normal Galaxies; Comastri et al. 2002). The common meaning of these definitions is that they lack evidence of accretion-driven activity in their optical spectra, in contrast with ``normal" Seyfert galaxies and quasars. Their X-ray luminosities ($\approx 10^{42}-10^{43}$ \lum), X-ray spectral shape and X-ray-to-optical flux ratio (X/O\footnote{Where X/O is defined as $X/O=\log{\frac{f_X}{f_R}}=\log{f_X}+C+\frac{R}{2.5}$.}$\sim -1$) suggest AGN activity of moderate strength. Originally discovered in early {\it Einstein} observations (Elvis et al. 1981) and named optically dull galaxies, the interest on the nature of these sources has gained a renewed attention after the discovery of several examples in XMM-{\it Newton} and {\it Chandra} surveys (Fiore et al. 2000; Comastri et al. 2002a,b; Georgantopoulos et al. 2005; Kim et al. 2006). Several possibilities were proposed in the literature in order to explain why a relatively luminous, hard X-ray source does not leave any significant signature of its presence in the form of emission lines. A simple explanation favoured by Moran et al. (2002) and more recently by Caccianiga et al. (2007) for faint sources in the \chandra\ deep fields and brighter object in the \xmm\ XBS survey, respectively, is dilution by the host galaxy starlight. The combination of optical faintness and lack of strong emission lines in the observed wavelength range for distant \chandra\ sources or the inadequate observing set-up among brighter nearby \xmm\ objects (Severgnini et al. 2003) may account for the \xbong\ properties. More in general, if the contrast between the host galaxy starlight and nuclear emission is high, AGN emission lines may easily be undetected. The physical reason may be ascribed to obscuration, most likely with a large covering factor, or to the fact that the lines are not efficiently produced by the central engine. If \xbongs\ are merely obscured AGN, two hypotheses may be envisaged: \par $\bullet$ In order to explain the multiwavelength properties of the \xbong\ prototype PKS~0312018, also known as P3, Comastri et al. (2002) suggested heavy obscuration by Compton-thick gas covering almost 4$\pi$ at the nuclear X-ray sources. In this way, no ionizing photons can escape to produce the narrow emission lines which are observed in ``normal" Type 2 narrow-line AGN which are thought to have a lower covering fraction following the AGN Unified Scheme (but see Section \ref{xray} for a recent re-analysis of the X-ray data of P3). \par $\bullet$ According to a detailed multiwavelength analysis of ``optically-dull" galaxies in the \chandra\ deep fields, Rigby et al. (2006) conclude that extranuclear dust in the host galaxy plays an important role in hiding the emission lines.\\ Alternatively, \xbongs\ may be members of a class, or classes, of {\it exotic} objects for which emission lines are either intrisically weak or absent: \par $\bullet$ Radiatively Inefficient Accretion Flows (RIAFs) are expected at accretion rates well below those inferred for Seyferts and quasars. A distinctive property of low accretion-rate flows is that the standard Shakura-Sunyaev accretion disk is truncated at a relatively large inner radius. As a consequence, it cannot generate the ``big blue bump" and enough UV photons to photoionize the line-emitting circumnuclear gas. The infalling gas is heated to high temperatures and emits a hard X-ray power-law by upscattering of low-energy seed photons. According to Yuan \& Narayan (2004), the SED of source P3 could be reproduced by a RIAF model. \par $\bullet$ \xbongs\ could be extreme BL Lac objects in which the featureless non-thermal continuum is much weaker than the host galaxy starlight. Following Fossati at al. (1998), \xbongs\ could belong to the low-luminosity tail of the blazar spectral sequence based on the anti-correlation between luminosity and frequency of the synchrotron peak. \par $\bullet$ A highly speculative possibility is that of a transient AGN phenomenon in the process of tidally disrupting a star. If this were the case, the X-ray emission should be witnessing the transient accretion phenomenon (see Komossa et al. 2004 for extreme variability events in ROSAT observations; see also Gezari et al. 2006 for a luminous flare observed in the GALEX Deep Imaging Survey). The transient is most likely over in subsequent follow-up optical observations. Finally, it should be noted that diffuse emission from a galaxy group, whose X-ray extended emission may have escaped detection in low signal-to-noise X-ray observations, is also possible and indeed observed in a few cases (Georgantopoulos et al. 2005). While a unique solution may not necessarily hold for all the \xbongs\ observed in different surveys, they represent a useful benchmark for a better understanding of the AGN activity and, as such, deserve further studies. Ideally, one would need sensitive, high-spatial resolution, multiwavelength observations from radio to X-rays. As a first step, in the following we use good-quality photometric near-infrared data obtained with ISAAC at VLT of four low-redshift ($z=0.1-0.3$) \xbongs, selected in the {\rm HELLAS2XMM} survey to search for the presence of a putative nucleus which has escaped detection in the optical spectroscopy. The rather obvious advantage of near-infrared data is that the effects of dust reddening are minimized and that the nuclear emission in this band should rise more rapidly than the stellar light due to the reprocessing by hot dust. At the same time, the excellent quality of the near-IR images makes possible to apply a surface brightness decomposition technique, already successfully applied by several authors (e.g. S\'anchez et al. 2004; Peng et al. 2006), to search for weak unresolved nuclear emission down to faint near-infrared magnitudes. We also discuss the broad-band properties of the four \xbongs\ using available multiwavelength data. Throughout the paper we assume a cosmological model with $H_0$ = 70 km s$^{-1}$ Mpc$^{-1}$, $\Omega_m$ = 0.3 and $\Omega_{\Lambda}= 0.7$. | We have presented a multiwavelength analysis of four \xbongs\ selected from the {\rm HELLAS2XMM} survey. For these sources, deep near-infrared images taken with ISAAC at VLT, good-quality optical spectra and \xmm\ data are available. Applying the morphological decomposition technique, we were able to detect the presence of a nuclear component in two out of the four sources (PKS~03120017 and PKS~03120018). There is no evidence of nuclear emission in the near-infrared in Abell~2690013; moreover, the X-ray appearance is consistent with an extended source. For source Abell~1835140, the near-infrared images reveal a complex morphology, where two sources are embedded in a common envelope. The main issue about the \xbong\ nature is whether they represent a truly distinct class or, rather, they are a mixed source population. Our results point towards the latter hypothesis. The results regarding the nature of the \xbongs\ with nuclear component and the lack of optical emission lines can be summarized as follows. \begin{itemize} \item Source PKS~03120017 and PKS~03120018 are well described by a mildly obscured ($E(B-V)= 0.5-0.8$) optically weak nucleus, responsible for the X-ray emission, hosted by a bright galaxy (mag$_{K_s}^{nucl}$--mag$_{K_s}^{host}$ $\sim 4$). \item The lack of optical emission lines cannot be attributed to observational limitations such as a inadequate observational setup or low signal-to-noise-ratio, at least for the objects in the present sample for which high-quality spectroscopic and photometric observations are available. \item We can safely discard for the two objects both a Compton-thick scenario (as found by Caccianiga et al. 2007 for a sample of elusive AGN in the XBSS) and an important blazar contribution. A RIAF solution seems to be supported by the estimated values of the Eddington ratios in the two objects and is not ruled out by the broad-band SED fitting, also because of the large number of free parameters in the RIAF model that can be tuned in order to reproduce the observed SEDs. A weak nuclear source, described by a standard accretion disk solution, but not powerful enough in the production of UV photons, would also provide an acceptable description of the observations. \item The presence of a thin nuclear gas and dust structure (as argued by Cocchia et al. 2007) covering 4$\pi$ at the nuclear source, combined with the low level of activity of the BH, could prevent the ionization of the narrow-line regions and produce also the extinction we measured.\\ The analysis of already obtained VLT/VIMOS-IFU (Integral Field Unit) observations could help in the detection of the typical AGN emission lines down to a very faint flux limit and in the study of the gas kinematics. \end{itemize} | 7 | 10 | 0710.1811 |
0710 | 0710.0671.txt | Astronomy provides a laboratory for extreme physics, a window into environments at extremes of distance, temperature and density that often can't be reproduced in Earth laboratories, or at least not right away. A surprising amount of the science we understand today started out as solutions to problems in astronomy. Some of this science was key in the development of many technologies which we enjoy today. This paper describes some of these connections between astronomy and technology and their history. | 7 | 10 | 0710.0671 |
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0710 | 0710.3163_arXiv.txt | This template file shows how to use the \texttt{aipproc} class to produce a paper with the correct layout for \emph{% AIP Conference Proceedings 8.5in x 11in double column}. A full description of the features supported by the \texttt{aipproc} class can be found in the \texttt{aipguide.pdf} document accompanying the distribution. Frequently asked questions can be found in the \texttt{FAQ.txt} document. | Infandum, regina, iubes renovare dolorem, Troianas ut opes et lamentabile regnum cruerint Danai; quaeque ipse miserrima vidi, et quorum pars magna fui. Quis talia fando Myrmidonum Dolopumve aut duri miles Ulixi temperet a lacrimis? Et iam nox umida caelo praecipitat, suadentque cadentia sidera somnos. | 7 | 10 | 0710.3163 |
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0710 | 0710.5486_arXiv.txt | { Type~Ia supernovae are believed to be white dwarfs disrupted by a thermonuclear explosion. Here we investigate the scenario in which a rather low-mass, carbon-oxygen (C~+~O) white dwarf accumulates helium on its surface in a sufficient amount for igniting a detonation in the helium shell before the Chandrasekhar mass is reached. In principle, this can happen on white dwarfs accreting from a non-degenerate companion or by merging a C~+~O white dwarf with a low-mass helium one. In this scenario, the helium detonation is thought to trigger a secondary detonation in the C~+~O core. It is therefore called the ``double-detonation sub-Chandrasekhar'' supernova model. By means of a set of numerical simulations, we investigate the robustness of this explosion mechanism for generic 1-\msol\ models and analyze its observable predictions. Also a resolution dependence in numerical simulations is analyzed. Hydrodynamic simulations of the double-detonation sub-Chandrasekhar scenario are conducted in two and three spatial dimensions. The propagation of thermonuclear detonation fronts, both in helium and in the carbon-oxygen mixture, is computed by means of both a level-set function and a simplified description for nuclear reactions. The decision whether a secondary detonation is triggered in the white dwarf's core or not is made based on criteria given in the literature. In a parameter study involving different initial flame geometries for He-shell masses of 0.2 and 0.1~\msol\ (and thus 0.8 and 0.9~\msol\ of C~+~O), we find that a secondary detonation ignition is a very robust process. Converging shock waves originating from the detonation in the He shell generate the conditions for a detonation near the center of the white dwarf in most of the cases considered. Finally, we follow the complete evolution of three selected models with 0.2~\msol\ of He through the C/O-detonation phase and obtain \nni-masses of about 0.40 to 0.45~\msol. Although we have not done a complete scan of the possible parameter space, our results show that sub-Chandrasekhar models are not good candidates for normal or sub-luminous type~Ia supernovae. The chemical composition of the ejecta features significant amounts of \nni\ in the outer layers at high expansion velocities, which is inconsistent with near-maximum spectra.} | \label{sec:int} One of the major uncertainties in modeling type~Ia supernovae (SN~Ia) originates from the unknown nature of the progenitor systems because neither observations nor theoretical models are yet conclusive \citep{Branch1995,Ruiz-Lapuente2000,Livio2000,Nomoto2003,Han2004,Napiwotzki2005,Stritzinger2006,Parthasarathy2007}. Over the past years, most studies of thermonuclear supernova explosions focused on models in which a thermonuclear flame is formed by a runaway near the center of a white dwarf (WD), composed of carbon and oxygen, once it has reached the Chandrasekhar mass ($M_\mathrm{Ch} \sim 1.4~\msol$) by accretion from a non-degenerate companion. Starting out in the sub-sonic deflagration mode of flame propagation, these models were shown to give rise to energetic explosions either in pure turbulence-boosted deflagrations \citep{Reinecke2002b,Gamezo2003,Roepke2005,Roepke2006a,Schmidt2006a,Schmidt2006b,Roepke2007c} or with a delayed triggering of a supersonic detonation phase \citep{Gamezo2004,Plewa2004,Roepke2007a}. There seems to be reasonable agreement of such models with the gross features of observed SNe~Ia \citep{Roepke2007c,Mazzali2007}. However, it remains unclear whether the Chandrasekhar mass model can account for the full range of SN~Ia observations. In particular, the mechanism of the sub-luminous events (\object{SN~1991bg} being a prototypical example) remains a puzzle \citep{Mazzali2007}. \citet{Stritzinger2006} even claim typical ejecta masses below $M_\mathrm{Ch}$ for a sample of well observed SNe~Ia, with the trend that the low-luminosity explosions eject less mass. Therefore, a long standing question is whether other progenitor channels contribute to the observed SN~Ia sample, and two alternatives have been suggested: the \emph{WD merger} or \emph{double degenerate scenario} (in contrast to single degenerate models, which consist of binaries with only one WD) and the \emph{sub-Chandrasekhar model}. In the present work we explore the latter. The basic idea of the sub-Chandrasekhar model is that if the accretion rate onto a WD is lower than about $1$ - $4 \cdot 10^{-8}~\msol \, \mathrm{yr}^{-1}$ the accreted He (or the H processed to He) is not fused steadily into C and O. Instead, after reaching a critical amount of He at relatively low densities, the He shell becomes unstable and detonates \citep[cf.][and references therein]{Woosley1986}. The detonation ignites most likely close to the bottom of the He shell, produces almost pure \nni, and can occur long before the WD reaches the Chandrasekhar limit. A second detonation may be triggered spontaneously when the He shell detonation shock wave hits the C~+~O core or with some delay after the shock has converged near the center of the WD (\emph{double detonation} models). This way, in the sub-Chandrasekhar mass models the conditions for thermonuclear runaway are caused by the compression in the shock wave and not by the high degree of degeneracy as in the Chandrasekhar mass models. Delayed double detonations have been studied extensively before. \citet{Woosley1994}, and \citet{Livne1990} carried out one-dimensional (1D) simulations and \citet{Livne1990a,Livne1991} considered two-dimensional (2D) setups. In 1D the He detonation ignites synchronously in a layer close to the bottom of the He shell and due to the spherical symmetry constraint the shock converges perfectly in the center. However, this is not a very realistic model. According to \citet{Livne1990} an ignition most likely happens in a single point leading to an off-center convergence of oblique shock waves in the core. Several other simulations predict successful directly ignited double detonations in two dimensions \citep{Livne1997,Arnett1997,Wiggins1997,Wiggins1998}. The smoothed particle hydrodynamics simulations by \citet{Benz1997} and \citet{Garcia-Senz1999} were carried out in 3D\@. \citet{Livne1990a} first reported an increased probability of the direct core ignition if the He detonation does not happen directly at the core--shell interface but at a certain distance above it. This way, the pressure jump can grow large enough before hitting the core. In this work, the results of a (restricted) parameter study are presented investigating the possibility of triggering the second detonation in two and three dimensions. In all our models the WD has a total mass of 1~\msol, but the (C~+~O)-core and the He-shell masses differ. We compute sequences of models with very different initial flame geometries in order to test the robustness of this explosion mechanism. Since this latter question is the main focus of our paper, most of the simulations are stopped once the conditions for a detonation are matched. However, for a few successful cases the energetics and nucleosynthesis of complete double detonations are computed. In the following chapter the model details will be described. Section~\ref{sec:sim} shows the simulation results, which are summarized and discussed in the last part of the paper. | \label{sec:con} We have studied the sub-Chandrasekhar model of SNe~Ia by means of a series of two-dimensional and a few three-dimensional simulations with different initial conditions. The numerical scheme used is based on the PPM algorithm, and the propagation of the detonation front is modeled applying the level set technique. This novel implementation allowed for an accurate treatment of the hydrodynamic features of the sub-Chandrasekhar model (such as shock waves and thin detonation fronts) in multiple dimensions. With our generic 1-\msol\ model we performed a parameter study involving different He masses and ignition geometries of the initial He detonation. We find that a detonation in the He shell can clearly trigger a second detonation in the core and at least for the mass configurations studied, the double detonation seems to be a very robust process, which works without any ``fine-tuning'' of our model parameters. In almost all of the simulations performed, the ignition conditions for a core detonation were reached at or near the center of the WD as a result of the convergence of the shock from the He shell detonation and in the few other cases the failure is likely a result of the finite numerical resolution. The high maximum densities and temperatures that are needed for detonation ignition and the fact that the shock is not diluted too much when propagating inwards against the pressure gradient are made possible through a geometrical shock amplification effect that appears in spherical symmetry: As the surface of the shock decreases its strength has to increase. The high compression that can be achieved theoretically in this way, especially if full convergence is reached and the shock is reflected (cf.\ Sect.~\ref{sec:res}), could even allow successful double detonations with considerably smaller He masses. The shock amplification that is reached, however, turns out to be resolution dependent on the numerical grid as it is coupled to the smallest resolvable shock surface. Taking this into account it is very likely that all our models would explode, if they were simulated with sufficiently high resolution. The question of whether an incineration could also happen directly at the core--shell interface, was not addressed here. This could of course prevent the spherical shock convergence mechanism from playing a role, at least in parts of the parameter space. Therefore, edge-lit detonations are postponed to a separate study. The complete double detonation simulations that were performed for the models z4.24A and s4.10A resulted in \nni-masses of about $0.40$ to $0.45~\msol$. Are the studied events thus good candidates for normal SNe~Ia? Most likely not. Starting with a He-shell detonation, all of our models show a significant amount of rapidly expanding \nni\ in the outer layers. In the observed spectra of normal and sub-luminous events, however, this has never been observed (cf.\ \citet{Branch1982,Branch1984a,Woosley1986,Arnett1997} and also the discussion in \citet{Livne1995}). The only exception is the super-luminous \object{SN~1991T}. Thus our solar-mass models most likely will not be able to explain normal or sub-luminous SNe~Ia and, given the robustness of the explosion mechanism of the model, the only conceivable explanation for the lack of observations of corresponding SN~Ia events is that the progenitors are not realized in nature -- or are very rare. For a better agreement with observations a reduction of the He-shell mass would be the most obvious choice. However, then a significantly larger core mass would most likely be required, because otherwise the He would not detonate. In this case the core density would be much higher resulting in a very large Ni mass. This would again not be a candidate for normal SNe~Ia, but it might be a promising candidate for objects like \object{SN~1991T}. Further work could cover model improvements like a more realistic treatment of nuclear reactions including the calculation of real reaction rates depending on the actual thermodynamic state of the burnt matter. This would make an investigation of the onset and the explosion dynamics of the core detonation possible. Also, more extended parameter studies, especially towards the lower end of possible He masses, would be interesting. For a comparison with observations the calculation of synthetic light curves and spectra would also be desirable. | 7 | 10 | 0710.5486 |
0710 | 0710.5023_arXiv.txt | {} {The principal aim of this project is to determine the structural parameters of the rapidly pulsating subdwarf B star PG 0911+456 from asteroseismology. Our work forms part of an ongoing programme to constrain the internal characteristics of hot B subdwarfs with the long-term goal of differentiating between the various formation scenarios proposed for these objects. So far, a detailed asteroseismic interpretation has been carried out for 6 such pulsators, with apparent success. First comparisons with evolutionary theory look promising, however it is clear that more targets are needed for meaningful statistics to be derived.} {The observational pulsation periods of PG 0911+456 were extracted from rapid time-series photometry using standard Fourier analysis techniques. Supplemented by spectroscopic estimates of the star's mean atmospheric parameters, they were used as a basis for the "forward modelling" approach in asteroseismology. The latter culminates in the identification of one or more "optimal" models that can accurately reproduce the observed period spectrum. This naturally leads to an identification of the oscillations detected in terms of degree $\ell$ and radial order $k$, and infers the structural parameters of the target.} {The high S/N low- and medium resolution spectroscopy obtained led to a refinement of the atmospheric parameters for PG 0911+456, the derived values being $T_{\rm eff}$=31,940$\pm$220 K, $\log{g}$=5.767$\pm$0.029, and $\log{\rm He/H}$=$-$2.548$\pm$0.058. From the photometry it was possible to extract 7 independent pulsation periods in the 150$-$200 s range with amplitudes between 0.05 and 0.8 \% of the star's mean brightness. There was no indication of fine frequency splitting over the $\sim$68-day time baseline, suggesting a very slow rotation rate. An asteroseismic search of parameter space identified several models that matched the observed properties of PG 0911+456 well, one of which was isolated as the "optimal" model on the basis of spectroscopic and mode identification considerations. All the observed pulsations are identified with low-order acoustic modes with degree indices $\ell$=0,1,2 and 4, and match the computed periods with a dispersion of only $\sim$ 0.26 \%, typical of the asteroseismological studies carried out to date for this type of star. The inferred structural parameters of PG 0911+456 are $T_{\rm eff}$=31,940$\pm$220 K (from spectroscopy) , $\log{g}$=5.777$\pm$0.002, $M_{\ast}/M_{\odot}$=0.39$\pm$0.01, $\log{M_{env}/M_{\ast}}$=$-$4.69$\pm$0.07, $R/R_{\odot}$=0.133$\pm$0.001 and $L/L_{\odot}$=16.4$\pm$0.8. We also derive the absolute magnitude $M_V$=4.82$\pm$0.04 and a distance $d$=930.3$\pm$27.4 pc.} {} | Subdwarf B (sdB) stars are evolved extreme horizontal branch stars with atmospheric parameters in the 20,000 K $\lesssim T_{\rm eff} \lesssim$ 40,000 K and 5.0 $\lesssim \log{g} \lesssim$ 6.2 range \citep{heber1986}. They are believed to be composed of helium-burning cores surrounded by thin hydrogen-rich envelopes and are characterised by a narrow mass distribution strongly peaked at $\sim$ 0.48 $M_{\odot}$ \citep{dorman1993}. While it is generally accepted that they evolved from the red giant branch (RGB) and constitute the immediate progenitors of low-mass white dwarfs \citep{bergeron1994}, the details of their formation are not yet understood. It does however seem clear that sdB progenitors lost a significant fraction of their envelope mass near the tip of the first RGB, leaving them with insufficient fuel to ascend the asymptotic giant branch (AGB) after core helium exhaustion. Plausible formation channels were modelled in detail by, e.g., \citet{han2002,han2003} and include binary evolution via a common envelope (CE) phase, stable Roche lobe overflow (RLOF), and the merger of two helium white dwarfs. These distinct evolutionary scenarios should leave a clear imprint not only on the binary distribution of sdB stars (CE evolution will produce sdB's in very close binary systems, RLOF gives rise to much longer period binaries, and the white dwarf merger results in single sdB stars), but also on their mass and hydrogen-envelope thickness distribution. Observational surveys focusing on radial velocity variations and the spectroscopic detection of companions have recently established that at least half of the sdB population reside in binaries (e.g. \citet{allard1994,ulla1998}), a significant fraction of them having short orbital periods from hours to days (\citet{green1997,maxted2001}). Accurate determinations of the stars' internal parameters on the other hand are harder to come by using traditional techniques; the mass has so far been measured only for the very rare case of an eclipsing binary \citep{wood1999}, while the envelope thickness eludes direct study. Fortunately, a small fraction ($\sim$ 5 $\%$) of sdB stars have been discovered to exhibit rapid, multi-periodic luminosity variations on a time-scale of hundreds of seconds, thus opening up the possibility of using asteroseismology to constrain their internal parameters. Since the near-simultaneous theoretical prediction \citep{charp1996, charp1997} and observational discovery \citep{kilkenny1997} of the so-called EC 14026 stars, both the modelling and measurement of their pulsational properties have come a long way (see \citet{fontaine2006} for a review). Simulating the pulsation spectra of a large grid of sdB models in terms of low-degree, low-order $p$-modes and numerically determining the "optimal" model that best fits a series of observed periodicities has so far resulted in the asteroseismological determination of the internal parameters for six EC 14026 pulsators: PG 0014+067 \citep{brassard2001}, PG 1047+003 \citep{charp2003}, PG 1219+534 \citep{charp2005a}, Feige 48 \citep{charp2005b}, EC 20117-4014 \citep{randall2006c}, and PG 1325+101\citep{charp2006}. These first asteroseismic results show promising trends as far as matching the expected mass and hydrogen shell thickness distribution is concerned, however more targets are needed to start assessing the importance and accuracy of the proposed formation channels. Here we present an asteroseismological analysis of the subdwarf B star PG 0911+456 based on photometry obtained with the new Mont4kccd at Mt. Bigelow, Arizona. In the next sections we describe the observations and frequency analysis, followed by the asteroseismic exploitation of the target and a discussion of the internal parameters inferred. | We obtained 57 hours of broad-band time-series photometry as well as high S/N low- and medium resolution time-averaged spectroscopy for the EC 14026 pulsator PG 0911+456. Our observations led to refined estimates of the star's atmospheric parameters and the detection of 7 independent harmonic oscillations, 4 more than were known previously. There was no sign of frequency splitting over the 68-day period during which the photometry was obtained, indicating a slow rotation rate. Fixing the effective temperature to the spectroscopic value and conducting an asteroseismic search in 3-dimensional $\log{g}-M_{\ast}-\log{q(\rm H)}$ parameter space enabled the identification of several families of models that could reproduce the observed periods to within less than 1 \%. While some of these were rejected from the outset due to obvious inconsistencies with the spectroscopic estimate of $\log{g}$ or implausible associated mode identifications, we retained two promising solutions for closer inspection. Unlike in some previous asteroseismological studies, it was not immediately obvious which model was to be preferred on the basis of the structural parameters alone; instead we used the inferred mode identification, in particular the degree index $\ell$, to discriminate between the two. The main difference between the solutions was that one identified a relatively high amplitude peak with an $\ell$=3 mode, while the other required only modes with $\ell$=0,1,2 and 4. Since detailed computations reveal $\ell$=3 modes to have extremely small disk-integrated amplitudes that would most likely not be detectable, we favoured the latter and adopted it as the optimal model. The inferred structural parameters for PG 0911+456 include the total stellar mass and the thickness of the hydrogen-rich shell, two quantities that can normally not be derived using other means but are invaluable for a detailed understanding of subdwarf B stars' evolutionary history. The total mass determined is smaller than that found for any EC 14026 star to date and places our target at the low-mass end of the predicted distribution. Similarly, the hydrogen envelope is measured to be thinner than that of most sdB's studied so far. If PG 0911+456 is confirmed to be a single star, as is suspected from its slow rotation, negligible radial velocity variation, and absence of a companion's spectroscopic or near-IR photometric signature, it may be the product of a WD merger according to the evolutionary channels proposed by \citet{han2002,han2003}. In this case, we would indeed expect the hydrogen envelope mass to be smaller than for an sdB having undergone a CE or RLOF phase. The low total mass derived would tend to support a non-canonical evolutionary history, even if it lies slightly below the mass distribution predicted from a WD merger. An obvious follow-up study for PG 0911+456 is to verify the asteroseismological solution found on the basis of additional observations. These could aim for a higher S/N level, thus enabling the detection of further pulsations and strengthening the constraints on the asteroseismic model. One of the main challenges we faced in the search of parameter space was the relatively large number of models that could account for the oscillations observed quite accurately; it would be very instructive to see whether any newly found periods can also be fit by the optimal model isolated. Observations containing wavelength-dependent information from which modes may be partially identified are another option. Following recent theoretical investigations into the amplitude-wavelength dependence of a mode on its degree $\ell$ (\citet{ramachandran2004}, \citet{randall2005}), there has been a surge in observational efforts to obtain multi-colour photometry of EC 14026 stars, most notably using the 3-channel CCD ULTRACAM (e.g. \citet{jeffery2005}) and the Mont4kccd predecessor, LaPoune I (e.g. \citet{fontaine2006}). While discriminating between low-degree modes with $\ell$=0,1,2 has proved extremely challenging, $\ell$=4 modes exhibit a more clearly distinguishable amplitude-wavelength behaviour (see e.g. Figure 26 of \citet{randall2005}). Given the necessary data, it should therefore be possible to confirm the identification of the two $\ell$=4 modes inferred for PG 0911+456 and thus verify the structural parameters computed. The work presented here constitutes the 7th detailed asteroseismological analysis of a rapidly pulsating subdwarf B star. While we estimate that the structural parameters of around 20 targets are required to start detailed comparisons with evolutionary theory, first tentative efforts in this direction look promising. Nevertheless, it is clear that there is still ample room for improvement on the modelling front. Firstly, the fact that the dispersion between the observed and theoretical periods of the optimal model is generally an order of magnitude higher than the measurement accuracy indicates remaining shortcomings in the models. We are currently working on full evolutionary (rather than envelope) "third-generation" models to address this problem. The reliability of the "optimal" models identified during the search of parameter space is another issue. Although we do apply cross-checks such as compatibility with non-adiabatic theory and spectroscopic values of the atmospheric parameters, the latter (especially $T_{\rm eff}$) must often be used to discriminate between, or constrain, regions of minimum $S^2$ and can no longer be employed as independent estimates. Moreover, the period ranges of unstable oscillations computed from our non-adiabatic pulsation code are sensitive mostly to $T_{\rm eff}$ and $\log{g}$ and are largely independent of the model mass and envelope thickness. It is therefore vital that additional checks are carried out with regard to the robustness of the "forward" approach if the structural parameters inferred from asteroseismology are to be compared with evolutionary predictions in a quantitative manner. The most obvious way of doing this is by detecting more frequencies from higher S/N observations or constraining the identification of the degree $\ell$ of individual modes from multi-wavelength time-series data. Such efforts are ongoing, and will likely prove invaluable for the future of sdB star asteroseismology. | 7 | 10 | 0710.5023 |
0710 | 0710.2435_arXiv.txt | Magnetic fields are observed to be associated with most structures in the universe. Observations indicate magnetic fields on stellar upto supergalactic scales. The field strengths vary from a few $\mu$G on galactic scale, upto $10^3$ G for solar type stars and upto $10^{13}$ G for neutron stars. Furthermore, the magnetic field structure depends on the object it is associated with. Thus, e.g., magnetic fields observed in elliptical galaxies show a different structure from those associated with spiral galaxies \cite{mag}. Magnetic fields in stars can be explained by the formation of protostars out of condensed interstellar matter which was pervaded by a pre-existing large scale magnetic field (see, e.g., \cite{rees}). An open problem remains to explain the origin of such large scale magnetic fields. There are different types of proposals. Ranging from processes on small scales, such as vortical perturbations and phase transitions to models taking advantage of the possibility of amplifying perturbations in the electromagnetic field during inflation in the early universe (see, e.g., \cite{rev-mag}). Inflation provides a mechanism to amplify perturbations in some field to appreciable size. In order for this mechanism to lead to primordial magnetic seed fields of cosmologically interesting strength, the corresponding lagrangian should not be conformally invariant. The Maxwell lagrangian describing linear electrodynamics is conformally invariant. There have been already a multitude of proposals to break the conformal invariance of the Maxwell theory \cite{tw}, e.g. by coupling to a scalar field \cite{mag-sc}, breaking Lorentz invariance \cite{mag-lor}, adding extra dimensions \cite{mag-ex} or a coupling to curvature terms \cite{mag-curv}. Here nonlinear electrodynamics is considered. It has its origins in the search for a classical singularity-free theory of the electron by Born and Infeld \cite{bi}. Later on it was realized that virtual electron pair creation induces a self-coupling of the electromagnetic field. For slowly varying, but arbitrarily strong electromagnetic fields the self-interaction energy was computed by Heisenberg and Euler (cf. \cite{he}-\cite{bb}). The propagation of a photon in an external electromagnetic field can be described effectively by the Heisenberg-Euler langrangian. Moreover, the transition amplitude for photon splitting in quantum electrodynamics is nonvanishing in this case. In principle, this might lead to observational effects, e.g., on the electromagnetic radiation coming from neutron stars which are known to have strong magnetic fields. \cite{bb,p-sp}. In particular, certain features in the spectra of pulsars can be explained by photon splitting \cite{puls}. Finally, Born-Infeld type actions also appear as a low energy effective action of open strings \cite{bi-strings,gh}. As was shown in \cite{dbi} the low energy dynamics of D-branes is described by the Dirac-Born-Infeld action. The model of the cosmological background that will be considered consists of a stage of de Sitter inflation followed by reheating and a standard radiation dominated stage. Quantum fluctuations in the electromagnetic field are excited within the horizon during inflation. Once outside the horizon they become classical perturbations. As mentioned above, in general, the conformal invariance of the four dimensional Maxwell field has to be broken in order to amplify the perturbations in the electromagnetic field significantly. Thus, here electrodynamics is considered to be nonlinear during the de Sitter stage. This could be motivated by the presence of possible quantum corrections to quantum electrodynamics at high energies. However, once inflation ends electrodynamics is described by standard Maxwell electrodynamics. Thus the subsequent evolution described by the standard model of cosmology is unchanged. | Observations of magnetic fields on large scales provide an intriguing problem. A possible class of mechanisms to create such fields is provided by inflationary models. Fluctuations in the electromagnetic field are amplified during inflation and provide a seed magnetic field at the time of structure formation which might be further amplified by a dynamo process. In general a sufficiently strong initial field strength can only be achieved if the conformal invariance of electrodynamics is broken. This has been realized, for example, in models where the Maxwell lagrangian has been coupled to a scalar field, to curvature terms, etc. or by breaking Lorentz invariance or adding extra dimensions. Here nonlinear electrodynamics has been considered. It has been assumed that whereas during the early universe electrodynamics is nonlinear it becomes linear at the end of inflation. In particular the evolution of the magnetic energy density has been studied in a model of nonlinear electrodynamics which is described by a lagrangian of the form $L\sim -\left[\left(F_{\mu\nu}F^{\mu\nu}\right)^2/\Lambda^8\right]^{\frac{\delta-1}{2}} F_{\mu\nu}F^{\mu\nu}$, where $\Lambda$ and $\delta$ are parameters. Originally the nonabelian version of this model had been proposed to describe low energy QCD \cite{pt}. Here this model has been chosen as it is a strongly nonlinear theory of electrodynamics which allows to study in a semi-analytical way the possible creation and amplification of a primordial magnetic field during de Sitter inflation. This is so since on the one hand the lagrangian only depends on one of the electromagnetic invariants, namely $X=\frac{1}{4}F_{\mu\nu}F^{\mu\nu}$, which leads to a significant simplification of the equations. On the other hand the power-law structure of the lagrangian make the equations simpler. Approximate solutions have been found in three regimes of approximation which describe the evolution of the ratio of the energy densities of the electric and magnetic fields during inflation. It is assumed that initially the energy density of the electric and magnetic field are of the same order. Furthermore, these initial fields are due to quantum fluctuations in the electromagnetic field during inflation. Whereas in the radiation dominated era, the energy density in the magnetic field decreases as $a^{-4}$, the electric field strength rapidly decays in the highly conducting plasma (see, e.g., \cite{tw,d93}). Solutions in closed form have been found and the resulting primordial magnetic field estimated. It has been shown that depending on the regime of approximation and the value of the Pagels-Tomboulis parameter $\delta$ primordial magnetic fields can be generated that are strong enough to seed a galactic dynamo. Thus we have provided an example of a theory of nonlinear electrodynamics where the nonlinearities act in a way as to amplify sufficiently an initial magnetic field. The energy-momentum tensor of the electromagnetic field can be cast in the form of an imperfect fluid. This has been found explicitly for the particular model of nonlinear electrodynamics under consideration here. Moreover, this allows to find the expression for the energy density $\rho$ of the fluid. Requiring that $\rho$ should be positive provides the bound $\delta\geq\frac{1}{2}$. In \cite{gfc} the possible creation and amplification of magnetic fields was studied in an inflationary model coupled to a pseudo Goldstone boson (see also \cite{tw}). In this case the lagrangian has the form $L\sim\frac{1}{2}\partial_{\mu}\theta\partial^{\mu}\theta -X+g_a\theta Y$, where $\theta$ is the axion field. This provides an example of a more general lagrangian having also an explicit dependence on $Y=\frac{1}{4}F_{\mu\nu}\;^{*}F^{\mu\nu}$. However, as it turns out the resulting primordial magnetic field is not strong enough in order to seed, for example, a galactic dynamo. In \cite{as} the model of \cite{gfc} was generalized to N axions. In this case it was found that at least the weaker bound of $r>10^{-57}$ can be satisfied. Here, in this work the creation of primordial magnetic fields in a particular model of nonlinear electrodynamics has been studied. It might be interesting to generalize this to lagrangians depending on both electromagnetic invariants $X$ and $Y$. | 7 | 10 | 0710.2435 |
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0710 | 0710.1600_arXiv.txt | Prior to the incineration of a white dwarf (WD) that makes a Type Ia supernova (SN Ia), the star ``simmers'' for $\sim 1000$ years in a convecting, carbon burning region. We have found that weak interactions during this time increase the neutron excess by an amount that depends on the total quantity of carbon burned prior to the explosion. This contribution is in addition to the metallicity ($Z$) dependent neutronization through the $^{22}$Ne abundance (as studied by Timmes, Brown, \& Truran). The main consequence is that we expect a ``floor'' to the level of neutronization that dominates over the metallicity contribution when $Z/Z_\odot\lesssim2/3$, and it can be important for even larger metallicities if substantial energy is lost to neutrinos via the convective Urca process. This would mask any correlations between SN Ia properties and galactic environments at low metallicities. In addition, we show that recent observations of the dependences of SNe Ia on galactic environments make it clear that metallicity alone cannot provide for the full observed diversity of events. | \label{sec:introduction} The use of Type Ia supernovae (SNe Ia) as cosmological distance indicators has intensified the need to understand white dwarf (WD) explosions. Of particular importance is the origin of the Phillips relation \citep{phi99}, an essential luminosity calibrator. Recent models demonstrate that it can be explained by large variations in the abundance of stable iron group elements \citep{woo07} with the dominant cause for diversity likely residing in the explosion mechanism \citep{maz07}. One additional variable is the metallicity of the WD core, which yields excess neutrons relative to protons due to the isotope $^{22}$Ne. This is usually expressed as \be Y_e = \sum_i \frac{Z_i}{A_i}X_i, \ee where $A_i$ and $Z_i$ are the nucleon number and charge of species $i$ with mass fraction $X_i$. The neutronization is critical for setting the production of the neutron-rich isotopes \citep{thi86}. If no weak interactions occur during the explosion, the mass fraction of $^{56}$Ni produced is simply $X(^{56}{\rm Ni}) = 58Y_e-28$, assuming $^{56}$Ni and $^{58}$Ni are the only burning products \citep{tbt03}. The neutronization also affects the explosive burning, including the laminar flame speed \citep{cha07}. However, the metallicity range of progenitors is not large enough to account for the full SNe Ia diversity (see \S 4), making it critical to explore all factors that determine $Y_e$. A potential neutronization site is the convective carbon burning core that is active for $\sim1000\ {\rm years}$ prior to the explosion. The hydrostatic evolution associated with this simmering phase terminates when the core temperature is sufficiently high that burning is dynamical \citep{nom84,ww86,woo04,ww04,kwg06}, and a flame commences \citep{tw92}. We show here that protons from the $^{12}$C($^{12}$C,$p$)$^{23}$Na reaction during simmering capture on $^{12}$C, and that subsequent electron captures on $^{13}$N and $^{23}$Na decrease $Y_e$. This process continues until either the explosion occurs or sufficient heavy elements have been synthesized that they capture the protons instead. In \S \ref{sec:rates}, we present simmering WD core models and summarize the reaction chains that alter $Y_e$. We find that one proton is converted to a neutron for each six $^{12}$C nuclei consumed for burning at $\rho<1.7\times 10^9\ {\rm g \ cm^{-3}}$. At densities above this, an additional conversion occurs from an electron capture on $^{23}$Na. Hence, the $Y_e$ in the core depends on the amount of carbon burned during simmering and the density at which it occurs, which we quantify in \S 3. We find that neutronization during simmering dominates for metallicities $Z/Z_\odot\lesssim2/3$. We conclude in \S \ref{sec:conclusion} by discussing the observations and noting where future work is needed. | \label{sec:conclusion} We have found that significant neutronization of the WD core occurs throughout the simmering stage of carbon burning until the onset of the explosion. If substantial energy is lost to the convective Urca process \citep[][and references therein]{les05}, then the neutronization is truncated by proton captures onto freshly synthesized heavy elements (resulting in eq. [\ref{eq:yemax}]). The main consequence is a uniform ``floor'' to the amount of neutronization that dominates over the metallicity dependent contribution for all progenitors with $Z/Z_\odot\lesssim2/3$. Given the likely significance this has for SNe Ia, more work needs to be done. In particular, at high ignition densities, heavy element electron captures and a full reaction network are needed to follow the resulting diverse collection of elements (see the discussion in \S 2.2). The convective Urca process % is another complication we have not addressed. In principle, if more energy is lost to neutrinos then more burning (and thus more neutronization) is required to make the burning dynamical. Assessing this will necessitate coupling a full nuclear network \citep{cha07b} to convective calculations. Such models would accurately determine $\eta$ rather than simply setting it to 1 or 2. In closing, we highlight some important features exhibited by recent observations of SNe Ia. It is clear that the amount of $^{56}$Ni produced in SNe Ia has a dynamic range ($0.1-1M_\odot$) larger than can be explained by metallicity or simmering neutronization. However, since an intriguing trend is the prevalence of $^{56}$Ni rich events in star-forming regions it is interesting to quantitatively explore how large the observed discrepancy is. Using the SNLS sample of Sullivan et al. (2006), Howell et al. (2007) found that the average stretch is $s=0.95$ in passive galaxies (e.g. E/S0's) and $s=1.05$ in star-forming galaxies. Using Jha et al's (2006) $\Delta M_{15}(B)-s$ relation and Mazzali et al.'s (2007) relation between $\Delta M_{15}(B)$ and $^{56}$Ni mass we get $0.58M_\odot$ ($s=0.95$) and $0.72M_\odot$ ($s=1.05$). Hence, amongst the large diversity, there is a tendency for SNe in star-forming galaxies to produce $\approx 0.13M_\odot$ more $^{56}$Ni than those in large ellipticals. Since the SN Ia rate scales with mass in ellipticals and star formation rate in spirals (Mannucci et al. 2005; Scannapieco \& Bildsten 2005; Sullivan et al. 2006), SNe from passive galaxies in the SNLS survey are from more massive galaxies than the SNe in star-forming galaxies (Sullivan et al. 2006). Using the mass-metallicity relation of Tremonti et al. (2004), our integration of the separate samples in Sullivan et al. (2006) yield average $12+\log({\rm O/H})=8.87$ in active galaxies and $9.1$ in ellipticals (solar value is $12+\log({\rm O/H})=8.69$). Due to the increase of ${\rm N/O}$ at high metallicities (Liang et al. 2006), the SNe in ellipticals have twice as much $^{22}$Ne content as those in spirals. From the relation of Timmes et al. (2003), this implies $\approx 5\%$ less $^{56}$Ni, whereas the observed decrement is $>15\%$. Explaining the observed decrement would require a contrast of $\Delta X(^{22}$Ne$)\approx 0.06$, or nearly 3 times solar. Although we have found that simmering enhances neutronization, the effect is not great enough ($\Delta Y_{e,\rm max}$ would give the same change in neutronization as doubling a solar metallicity), and a diverse set of core conditions would still be required. A large enhancement could be present in the core if substantial gravitational separation had occurred \citep{bh01,db02}, yet convective mixing during simmering will reduce it based on the fraction of the star that is convective. For a convection zone that extends out to $M_\odot$, the resulting $^{22}$Ne enhancement would be at most $\approx 30\%$. | 7 | 10 | 0710.1600 |
0710 | 0710.1752_arXiv.txt | The early evolution of dense stellar systems is governed by massive single star and binary evolution. Core collapse of dense massive star clusters can lead to the formation of very massive objects through stellar collisions ($M\geq$ 1000\,\msun). Stellar wind mass loss determines the evolution and final fate of these objects, and decides upon whether they form black holes (with stellar or intermediate mass) or explode as pair instability supernovae, leaving no remnant. We present a computationaly inexpensive evolutionary scheme for very massive stars that can readily be implemented in an N-body code. Using our new N-body code 'Youngbody' which includes a detailed treatment of massive stars as well as this new scheme for very massive stars, we discuss the formation of intermediate mass and stellar mass black holes in young starburst regions. A more detailed account of these results can be found in \cite{Belkusetal}. | 7 | 10 | 0710.1752 |
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0710 | 0710.4946_arXiv.txt | Thermal instability of partially ionized plasma is investigated by means of a linear perturbation analysis. According to the previous studies under the one fluid approach, the thermal instability is suppressed due to the magnetic pressure. However, the previous studies did not precisely consider the effect of the ion-neutral friction, since they did not treat the flow as two fluid which is composed of ions and neutrals. Then, we revisit the effect of the ion-neutral friction of the two fluid to the growth of the thermal instability. According to our study, the characteristic features of the instability are the following four points: (1) The instability which is characterized by the mean molecular weight of neutrals is suppressed via the ion-neutral friction only when the magnetic field and the friction are sufficiently strong. The suppression owing to the friction occurs even along the field line. If the magnetic field and the friction are not so strong, the instability is not stabilized. (2) The effect of the friction and the magnetic field is mainly reduction of the growth rate of the thermal instability of weakly ionized plasma. (3) The effect of friction does not affect the critical wavelength $\lambda _{\rm F}$ for the thermal instability. This yields that $\lambda _{\rm F}$ of the weakly ionized plasma is not enlarged even when the magnetic field exists. We insist that the thermal instability of the weakly ionized plasma in the magnetic field can grow up even at the small length scale where the instability under the assumption of the one fluid plasma can not grow owing to the stabilization by the magnetic field. (4) The wavelength of the maximum growth rate of the instability shifts shortward according to the decrement of the growth rate. This is because the friction is effective at rather larger scale. Therefore, smaller structures are expected to appear than those without the ion-neutral friction. Our results indicate the friction with the magnetic field affects the morphology and evolution of the interstellar matter. In summary, the ion-neutral friction is important for the evaluation of the thermal instability in weakly ionized plasma along and perpendicular to the magnetic field. | The recent progress of the interstellar medium (ISM) observations has established that the small and tiny scale structures of ISM is very ubiquitous. Historically, the initial observational target is the neutral phase of ISM. Indeed, the tiny scale structures are discovered by 21 cm absorption line observations against quasars with VLBI techniques \citep{Di76}. The results are confirmed by \citet{Dia89}. To look for the small scale structures in the cold neutral phase of the ISM furthermore, the observations of 21 cm absorption lines have been also performed against pulsars \citep{Fr94}. \citet{Me96} observe optical spectral lines to find small structures against close binary stars. The images of the small-scale H \textsc{i} have been taken with the MERLIN array \citep{Da96}. These results are established by higher angular resolution observations with VLBA and VLA toward quasars \citep{Fa98,FG01}, which show small clumps on the order of a few AU in neutral Galactic H \textsc{i} clouds. Cold H \textsc{i} clouds have significant structure in subparsec scales \citep{Gi00,Br05}. Such small and tiny scale structures has been detected in the local interstellar medium as H \textsc{i} absorption lines, although their column density is very small \citep{BK05}. In addition to the small and tiny scale structures of H {\rm \textsc{i}} cloud, the picture of the planetary nebula NGC 7293 shows fine small structures and knots \citep{Ro02,OD04}. Even in the starforming regions, there are variety small scale structures. For example, \citet{La95} observe some clumps with size from 0.007 to 0.021 pc in the Taurus Molecular Cloud 1 (TMC1), in particular, core D. The mass of these fragments is estimated to be $\lesssim 0.01 - 0.15$ $M_{\odot}$. SPITZER has begun producing higher spatial resolution mid-infrared maps \citep{Ch04}, and revealed the fine structure of the starforming region. We think that it is possible for proto-brown dwarfs of $< 0.08$ $M_{\odot}$ to exist there. A high resolution observation in future is expected to reveal hidden, small structures, as well as substellar objects with very small mass. In any ways, the tiny-scale structure seems to be ubiquitous, not associated with large extinction \citep{He97}. To study the origin of the small-scale structures of ISM is important to understand the evolution of and structure formation in ISM. Especially, in the starforming regions, the small, low-mass structures can relate to the low-mass cutoff of the initial mass function, the coagulation unit to form the massive stars, and so on. Then, we should investigate the physical origin of these small and tiny structures in the partially ionized medium, since the cold H {\rm \textsc{i}} and molecular clouds are weakly ionized. By the way, it is not easy for the small and tiny scale structure to form as a result of gravitational instability since their size is much smaller than the Jeans length. According to \citet{La95}, indeed, the small-scale structures appear to be gravitationally unbound. This suggests that some fragmentation mechanisms but pure Jeans gravitational instability may be important in the clouds. We expect the thermal instability as this mechanism, and revisit it in this paper. The basics of the thermal instability has been summarized in \citet{F65}. When the following condition is satisfied, the system is thermally unstable: If the cooling becomes efficient as the temperature decreases, the thermal energy gets lost more and more. If the cooling becomes efficient during the fluid contraction, the system shrinks on and on. Importantly, the critical length scale of the thermal instability is smaller than that of the dynamical instability like the Jeans instability. That is, even if a system is stable against the gravitational instability, the system can become thermally unstable. Then, the thermal instability can be the physical origin of the smaller-scale formation than the dynamical instability. Indeed, \citet{KI02} propose that the clumpiness in clouds emerges naturally from their formation through the thermal instability. Their two-dimensional calculations follow the fragmentation into small cloudlets that result from the thermal instability in a shock-compressed layer. \citet{BL00} investigate the cooling and fragmentation of optically thin gas with a power-law cooling function. According to them, small-scale perturbations have the potential to reach higher amplitude than large-scale fluctuations. Thermal instability can be important for the structure formation. We sketch ISM structure formation such as molecular clouds as follows. The hot ionized ISM cools to be cold and weakly ionized clouds. Ionization degree of ISM decreases as it evolves, and then the ISM fluid is often partially ionized. In partially ionized fluid, an ion component and neutral component interact each other, exchanging the momentum (i.e., ion-neutral friction). Especially, the partially ionized plasma in a magnetic field could not be treated as one fluid. This is because although the ion component in the partially ionized plasma is directly influenced by a magnetic field, the neutral component does not directly feel a magnetic field. In weakly ionized fluid, the neutral component is affected by a magnetic field via the ion-neutral friction. For example, ambipolar diffusion takes an important role in dynamical evolution owing to the gravitational instability of ISM into protostar \citep{MeSp56, Na76}. If the fluid frozen in a magnetic field contracts without ambipolar diffusion, the magnetic field becomes too large and suppresses the growth by the magnetic pressure and tension. In addition, \citet{Na79} points out the possibility that ISM with sufficient magnetic flux quasistatically evolves into protostar. He also notices that the time scale of plasma drift depends on the ionization degree. The relation among the ion-neutral friction, the magnetic field, and the amount of the ion component is important to understand the detailed process. We here emphasize that the MHD approximation is not always applicable for our purpose in this paper. The thermal instability is effective in small structure formation, during which a neutral component is not always enough frozen in an ion component dynamically to be treated as one fluid. There is a possibility that the ion component can drift away owing to ambipolar diffusion, and that the system can contract owing to thermal instability. Then, in the present paper, to study how the ion-neutral drag and magnetic field influence the growth of the thermal instability, we treat the plasma as the two fluid of ions and neutrals. We assume that the ionization degree is very small, because we are interested in the final stage of the formation of a molecular cloud, in which the ionization degree decreases very much. We focus on the condensation mode of thermal instability, since it is important in structure formation, rather than the oscillation (overstable) mode. In \S2, the problem is formulated. The property of the dispersion relation of the instability is presented in \S3. Several discussions relating to the instability are found in \S4. Applicability of our results to cold H {\rm \textsc{i}} medium is briefly examined in \S5, and then we summarize the paper in \S6. | \label{S_Discuss} We study the basic property of the thermal instability of the weakly ionized plasma in the previous sections. In this section, basing on our understanding the thermal instability, we shall discuss the structure formation. \subsection{Critical Wavelength and Thermal Conduction} \label{subsec:critical_wavelength_conduction} First in this section, the critical wavelength of the weakly ionized plasma in the magnetic field is discussed. As mentioned in section \ref{S_Results}, the critical wavelength of the mode relating $\mu_{\rm n}$ is not changed by the magnetic field with the friction as long as the thermal conduction is not changed by the magnetic field. In fact, the thermal conductivity is varied owing to the magnetic field. In this meaning, we may oversimplify the problem in this paper. However, the thermal conductivity will decrease as the magnetic field strength increases, and then the critical length become smaller, according to equation (\ref{eq:appendix:criticallength:neutral}). Thermal conduction erases the temperature perturbation and thermally stabilizes the system, especially at smaller scale. The larger thermal conduction makes the critical wavelength longer, owing to the criterion \ref{crite_cond} or \ref{crite_ion}, according to \citet{F65}. On the other hand, the magnetic field lessens the thermal conduction of the ion component, because the ion component winds around a magnetic field. Therefore, we can insist that the critical length for the thermal instability of the weakly ionized plasma can not be enlarged by the friction and the magnetic field at least, while the effect of the magnetic field can influence not only the ion but the neutral component via the ion-neutral friction. This statement has universal validity for the thermal instability of the weakly ionized plasma. The effect of the friction and the magnetic field is mainly reduction of the growth rate of the thermal instability. On the other hand, the magnetic field directly enlarges the critical wavelength of the mode characterized by $\mu_{\rm i}$ as expected in the one fluid plasma model. \subsection{Spatial Distribution of Magnetized Interstellar Medium} In the previous section, we learn that the suppression of the thermal instability is characterized by the ion-neutral friction, strength of the magnetic field, and the direction of the field. Interestingly, the suppression of the instability is different owing to the magnetic field direction. This suggests the anisotropic nature of the suppression of the instability is imprinted in the morphology of the structure. In this standing point, we shall discuss how the structure formation proceeds in the magnetized ISM. \subsubsection{Growth Rate Parallel to the Magnetic Field} In an usual one fluid approach, the motion of the plasma parallel to the magnetic field is regarded not to be affected. However, Figure \ref{B0C003} suggests that the growth rate of the thermal instability at larger scale decreases via the ion-neutral friction. This suppression of the instability yields the delay of the evolution in comparison to the evolution under the MHD approximation without the ion-neutral friction. \subsubsection{Distribution of Partially Ionized Plasma Perpendicular to the Magnetic Field Line} \label{sss_spread_of_part} Here, we suggest that the spatial distribution of the partially ionized plasma is elongated, whose semi-major axis is vertical to the magnetic field line. Figures \ref{B0C003} and \ref{B06C003}, both of which friction strengths are the same, show the feasibility. The thermal instability in Figure \ref{B0C003} corresponds to the mode along the magnetic field line since $\mbox{\boldmath$B$} = 0$, while that in Figure \ref{B06C003} is vertical to the magnetic field line. The growth of thermal instability vertical to the line is suppressed by both of the magnetic field and the friction, while the growth along the line is suppressed only by the friction. Thus, the growth rate along the magnetic field line is larger than that vertical to the line, and then the resultant morphology of the structure via the thermal instability is elongated. \subsubsection{Difference of Spatial Distribution of Ion and Neutral Components} \label{ssS_Difference_in_Spread} We insist the growth rate depending on the direction of the field results in the morphology of interstellar structure. Furthermore, we learn in \S3 the mode characterized by $\mu_{\rm i}$ are stabilized by the friction and field, while that by $\mu_{\rm n}$ can grow as found in Figure \ref{B06C003}. Figure \ref{B600C5} shows that although the mode characterized by $\mu_{\rm i}$ is completely stabilized, the mode characterized by $\mu_{\rm n}$ is still unstable. Consequently, we can insist that the neutral component condenses owing to the thermal instability, while the ion component still keeps spreading. It is noticed that the ion component may condense since it is dragged by the neutral component contracting owing to the thermal instability. However, once the distribution of the ion and the neutral components separates and the motion of the two becomes independent there, the ion component in the area is thermally stable owing to the magnetic field. This means the magnetic pressure lets the ion component be spread in the area even if the neutral component condenses. This thin ion component left behind by the contracting neutral component will be observed as a halo around H {\rm \textsc{i}} clouds, and it emits, especially some forbidden lines. This halo possibly has wide distribution vertical to the magnetic field, as discussed in subsubsection \ref{sss_spread_of_part}. This is because the ion component can condense, parallel to the magnetic field like the neutral. The magnetic pressure can not prevent the motion of the ion component efficiently as found in Figure \ref{B0C003}. Thus, the morphology of the ion halo is also elongated according to the orientation of the field line. \subsubsection{Origin of the Small Clumpinesses of Partially Ionized Plasma} \label{subsubsec:small_clumpinesses_of_partial} It is suggested that there are small clumpy structures of ISM. Here, we examine whether the clumpiness originates via the thermal instability. According to Figures \ref{B600C03} and \ref{B600C5}, the growth rate of the thermal instability of partially ionized plasma depends on the strength of the friction. The scale length at the maximum growth rate of the instability becomes smaller, owing to the suppression via the friction which is effective at larger scale. Thus, we can expect the small clumpiness appears selectively via the thermal instability since the ion-neutral friction suppresses the large scale condensations. If we are interested in the small clumpiness of ISM, we should mention about the thermal conduction in relation to the magnetic field. According to subsection \ref{subsec:critical_wavelength_conduction}, the small clumpiness may the results of the decrement of the conduction. It is noticed that the conduction along the field line is not so affected, and then there is still ambiguity for the reduced conduction coefficient to perform the key role for the origin of the small clumpiness. It is also noticed that the thermal instability is a slow process relative to the dynamical processes. If molecular clouds form from the H {\rm \textsc{i}} medium via the gravitational instability, the small clumpiness inside the molecular clouds appears after the formation of the clouds. This trend is enhanced if the effects of the magnetic field and friction are concerned. This is because the friction and the field suppress the large scale growth, while the short wavelength modes are still unstable. In a enough long span for the growth, the tiny structures can appear in ISM and affect the evolution of its structure especially as the non-linear effects. Thus, if we want to know how faint and small structures come to be observable, the growth time-scale of the thermal instability should be precisely examined. At this time, we never forget the effect of the magnetic field and the ion-neutral friction. \subsection{Ineffectiveness of the Magnetic Field} \label{S_Discuss_Satu} The modes in Figures \ref{B600C03} and \ref{B6000C03} have similar growth rate, while the magnetic fields are significantly different in each other. It implies that the suppression of the growth by increasing the magnetic field with the fixed friction has limits. By the way, a dispersion relation of only ionized fluid in a magnetic field, which is derived by equalizing the second bracket on the left-hand side in equation (\ref{dis-rela}) to be zero, is reduced to \begin{equation} n^3 + n^2 v_{\rm si} k_{\rm iT} = 0 , \end{equation} when $k$ approaches $0$. This is solved to be \begin{equation} n=0 ~ {\rm or}~ n = -v_{\rm si} k_{\rm iT} . \end{equation} Thus, we find that this solution is independent of a magnetic field. Even if the magnetic field becomes strong, the value of the growth rate at $k \sim 0$ is scarcely changed, where the friction is more effective than at small wavelength. Therefore, once the effect of the friction is fixed, increasing the magnetic field strength beyond a threshold level apparently becomes ineffective for the suppression of the growth. Not only information of the magnetic field but the friction is necessary to investigate the evolution of the thermal instability of the weakly ionized plasma. \subsection{Comment on One Fluid Approximation} The formulation of the fully one component plasma can not be always applicable to study the thermal instability of weekly ionized plasma. This is because there is the growing mode even when there exist the magnetic field and friction, whose growth rate is comparable to that of the neutral component when there are no magnetic field and friction. Indeed, this trend is found from the comparison among Figures \ref{B06C003}, \ref{B600C03} and \ref{B6000C03}. In addition, regarding weakly ionized plasma including the friction effect, the friction does not change the critical wavelength $\lambda_{\rm F}$. The friction tends to efficiently reduce the growth at much larger scale than the $\lambda_{\rm F}$. This feature yields that the critical wavelength in the weakly ionized plasma is not changed by the magnetic field via the friction, as mentioned in section \ref{S_Results} and subsection \ref{subsec:critical_wavelength_conduction}. Even if the magnetic field becomes strong, only the growth rate is reduced, keeping the critical wavelength constant. This is also the difference from the simple analysis of the one component magnetized plasma. The previous study of the one component plasma shows the stabilization at the smaller scales, while the current study insist $\lambda_{\rm F}$ is unchanged even if the effect of the field is concerned. \label{S_Summary} The thermal instability of the weakly ionized fluid is investigated with a linear perturbation analysis. The plasma is assumed to consist of two fluids of the ion and the neutral components. With this approach, the effect of the friction between the ion and the neutral components and the magnetic field are important. Here, the properties of the thermal instability of weakly ionized plasma and the observational implications are summarized. \begin{enumerate} \item Modes relating to $\mu_{\rm n}$ are not stabilized by the magnetic field. This is because the neutral component feels the field via the ion-neutral friction. \item Modes relating to $\mu_{\rm i}$ are directly stabilized by the magnetic field. This means that when the magnetic field is large, the ion component is hard to follow the motion of the neutrals which condensate owing to the thermal instability. \item The growth rate of the mode vertical to the magnetic field is reduced by the magnetic field and the ion-neutral friction. \item The growth rate of the mode along the magnetic field decreases owing to the ion-neutral friction, especially at large scale. The suppression of the instability is more ineffective than that of the mode vertical to the magnetic field. \item The critical wavelength for the thermal instability is not affected by the friction. The effect owing to the two fluid approximation is the reduction of the growth rate of the thermal instability. The magnetic field makes the critical wavelength of modes relating $\mu_{\rm i}$ larger. \item To study the thermal instability of the partially ionized plasma, one fluid approximation is not always useful. \item The ion-neutral friction and the magnetic field affect the distribution or morphology of ISM, especially after the long time compared to the free-fall time. The difference between the growth rate along and perpendicular to the magnetic field is important. The partially ionized plasma possibly is elongated perpendicular to the magnetic field. In addition, the neutral component and the ion component of weakly ionized fluid in the magnetic field are possibly separated each other. Then, the neutral component condenses owing to thermal instability, while the ion component left behind by the contracting neutral component still keeps spreading. Therefore, it is possibly observed that the weakly ionized fluid is elongated vertically to the magnetic field and surrounded by a halo which is rare and diffuse ions and emits forbidden lines. \end{enumerate} To conclude, the treating both of the independent motions of the neutral and the ion components yield the ion-neutral friction. The friction is important for the thermal instability of weakly ionized fluid, especially when studying the structure formation such as molecular cloud at its final stage of formation. Our study indicates that the fully ionized plasma approximation or totally neutral fluid approximation is not always applicable to weakly ionized plasma. | 7 | 10 | 0710.4946 |
0710 | 0710.1564_arXiv.txt | {% In a self-absorbed synchrotron source with power-law electrons, rapid inverse Compton cooling sets in when the brightness temperature of the source reaches $T_{\rm B}\sim10^{12}\,$K. However, brightness temperatures inferred from observations of intra-day variable sources (IDV) are well above the "Compton catastrophe" limit. This can be understood if the underlying electron distribution cuts off at low energy.} {% We examine the compatibility of the synchrotron and inverse Compton emission of an electron distribution with low-energy cut-off with that of IDV sources, using the observed spectral energy distribution of S5~0716+714 as an example.} {% We compute the synchrotron self-Compton (SSC) spectrum of monoenergetic electrons and compare it to the observed spectral energy distribution (SED) of S5~0716+714. The hard radio spectrum is well-fitted by this model, and the optical data can be accommodated by a power-law extension to the electron spectrum. We therefore examine the scenario of an injection of electrons, which is a double power law in energy, with a hard low-energy component that does not contribute to the synchrotron opacity.} {% We show that the double power-law injection model is in good agreement with the observed SED of S5~0716+714. For intrinsic variability, we find that a Doppler factor of $\doppler\geq30$ can explain the observed SED provided that low-frequency ($<32\,$GHz) emission originates from a larger region than the higher-frequency emission. To fit the entire spectrum, $\doppler\ge65$ is needed. We find the constraint imposed by induced Compton scattering at high $T_{\rm B}$ is insignificant in our model. } {% We confirm that electron distribution with a low-energy cut-off can explain the high brightness temperature in compact radio sources. We show that synchrotron spectrum from such distributions naturally accounts for the observed hard radio continuum with a softer optical component, without the need for an inhomogeneous source. The required low energy electron distribution is compatible with a relativistic Maxwellian.} | \label{intro} Observations of many extra-galactic radio sources have found rapid flux variations at radio frequency \citep[e.g.][]{kedziorachudczeretal01}, some of which fluctuate over a time scale of a day or less. They are referred to as intra-day variable sources (IDV). The variability time scale is often used to constrain the size of the source based on causality arguments. Using this constraint, one can derive a variability brightness temperature \citep{wagnerwitzel95} \eqb T_{\rm var}=4.5\times10^{10}F_{\nu}\left({\lambda d_{\rm L}\over t_{\rm obs} (1+z)}\right)^2\, {\rm K} \eqe where the flux density $F_\nu$, wavelength $\lambda$, luminosity distant $d_{\rm L}$, and observed variability time scale $t_{\rm obs}$ are measured in Jy, cm, Mpc, and days, respectively. The high radio flux frequently measured in IDV sources implies an extremely high brightness temperature, often many orders of magnitude above $10^{12}\,$K. \citet{kellermannpaulinytoth69} have shown that, assuming the electron distribution follows a single power law, the luminosity of the inverse Compton scattered photons exceeds that of the synchrotron photons when the brightness temperature of the source reaches $\sim10^{12}\,$K. Above this threshold, rapid cooling of the relativistic electrons due to inverse Compton scattering --- the \lq\lq Compton catastrophe\rq\rq\ --- forbids a further increase in the brightness temperature \citep[see e.g.][for a recent review of the brightness temperature problem]{kellermann02}. The limiting value is even lower, $T_{\rm B}<10^{11}\,$K, if the magnetic field and particle energy density of the source is driven towards equipartition \citep{readhead94}. The observed variability in some sources can be interpreted as the result of extrinsic effects, which, at first sight, relaxes the size constraint. For example, the flux variations of PKS~1519$-$273 and PKS~0405$-$385 are convincingly identified as interstellar scintillation. Nevertheless, all realistic models of the scintillation mechanism impose a new constraint on the size and require a brightness temperature of $T_{\rm B}>10^{13}\,$K in some cases \citep{macquartetal00,rickettetal02}, far exceeding the limit imposed by the Compton catastrophe. A prevalent feature associated with IDV sources is a flat or inverted spectrum ($\alpha\leq0$, with flux $F_\nu\propto\nu^{-\alpha}$) at radio-millimeter wavelengths \citep[e.g.,][]{gearetal94,kedziorachudczeretal01}. Optically thick synchrotron emission from power-law electrons rises as $\nu^{5/2}$, too fast to account for the observed spectra. Optically thin synchrotron emission in the scope of the conventional interpretation of the synchrotron theory has a flux $F_\nu\propto\nu^{-(s-1)/2}$, where $s$ is the power-law index of the electrons ($\diff\nelec/\diff\gamma\propto\gamma^{-s}$). If $\alpha=(s-1)/2\leq0$, the number density of electrons diverges towards high $\gamma$. Imposing a high-energy cut-off in the electron spectrum avoids the divergence and may account for the commonly observed spectral steepening at optical frequencies, but \citet{marscher77} showed that electron spectra with $s\leq1$ would result in a high flux between infrared and optical frequencies that is not supported by observations. The most common interpretation of the flat or inverted spectra is, therefore, a superposition of many synchrotron spectra within an inhomogeneous source \citep[e.g.][]{debruyn76,marscher77,blandfordkoenigl79}. In \cite{kirktsang06}, we discussed a synchrotron self-Compton model in which the electron distribution is monoenergetic. The lack of low-energy electrons enables more GHz photons to emerge from the source, allowing a higher brightness temperature to be observed without initiating catastrophic cooling. We found that a temperature of up to $T_{\rm B}\sim10^{14}\,$K at GHz frequencies is possible with only a moderate Doppler boosting factor of $\sim10$. In \citet{tsangkirk07}, we discussed the parameters of the monoenergetic model and showed that the assumption of equipartition of energy in the source does not prevent the Compton catastrophe. We also showed that an injection of highly relativistic electrons or strong acceleration in the source cannot produce temperatures much higher than our limit due to copious electron-positron pair production. In this paper, we examine the spectral properties of synchrotron emission from monoenergetic electrons and from an electron distribution that is a double power law in energy, by comparing the model spectra with the observations of S5~0716+714, a BL~Lac object that is one of the brightest known IDV sources, as well as a gamma-ray blazar \citep{hartmanetal99}. In doing so, we assume that the dominant targets for inverse Compton scattering are produced within the source (SSC model). The emission from gamma-ray blazars can also be interpreted in the context of models in which the target photons are created externally (EC model), for example in the broad line region, the accretion disk, or a molecular torus \citep{sokolovmarscher05}. However, in many sources there is no observational evidence of a significant external photon source. This is the case for S5~0716+714, where, despite much effort over the past three decades, no emission lines have been detected \citep[e.g.,][]{bychkova06}. Furthermore, XMM-Newton observations of S5~0716+714 in 2004 analysed by \citet{ferreroetal06} and \citet{foschinietal06} show two spectral components in the $0.5-10\,$keV band, whose variability properties appear to favour the SSC interpretation. The recent extensive simultaneous observations of this object from radio to optical frequencies by \citet{ostoreroetal06}, together with INTEGRAL pointings at GeV $\gamma$-ray energies during the same period, provide the best test for our model. In the following, we present the computation of the stationary electron distribution and the resulting synchrotron and inverse Compton spectra. The model spectra computed using the monoenergetic electron approximation, as described in \citet{tsangkirk07}, are presented first. Although adequate for the radio emission, the monoenergetic model cannot reproduce the entire spectrum of S5~0716+714. We therefore investigate an electron distribution that is a double power law in energy --- a hard low-energy part that softens to a high-energy tail above a characteristic energy. In this way, the inverted optically thin radio emission is retained and complemented by nonthermal synchrotron emission from the high energy tail. In section~\ref{parameter}, we briefly describe these injection models. The resulting stationary electron distribution is calculated in section~\ref{stationarysoln} and used for the computation of the synchrotron and inverse Compton spectra. In section~\ref{sed}, we compare the predictions of these models with the observed spectral energy distribution (SED) of the source to S5~0716+714. Our findings and some limitations of our approach are discussed in section~\ref{discussion} and our conclusions presented in section~\ref{conclusion}. | \label{conclusion} Using the specific case of S5~0716+714 as an example, we confirm that it is possible to produce high brightness temperatures at GHz frequencies in compact radio sources without the onset of catastrophic cooling, provided that the radiating particles have a distribution that is sufficiently hard below a characteristic. In addition, we show qualitatively that induced Compton scattering is insignificant in sources with a low-energy electron cut-off despite the high brightness temperature, the underlying reason being the low occupation number of the photons that can couple with the electrons at the cut-off energy. The model where an electron distribution that is a double power law in energy, peaking at $\gammap$, is injected into the source offers more flexibility at higher frequencies in the synchrotron spectrum (from infrared to optical) at the expense of more free parameters, compared to either monoenergetic or single power-law distributions. These parameters should be constrained by simultaneous observations due to the highly variable nature of IDV sources. In the case of S5~0716+714 where such data is available, the spectral break at about $230\,$GHz determines the value of $\gammap$, the optical data at $5\times10^{14}\,$Hz gives the lower limit of $\numax$, as well as constraining the spectral index $s_2$, and the INTEGRAL upper limits give the upper limit of $\numax$ and also constrain the value of $r_{\rm p}$, which in turn determines the electron density. The example of S5~0716+714 illustrates several important spectral properties of an electron distribution with a low-energy cut-off, as described in the previous sections. The most noticeable feature is the hard, inverted, optically thin synchrotron spectrum, spanning a wide frequency range, which is a prevalent feature in compact radio sources at radio frequencies \citep[e.g.,][]{gearetal94,kedziorachudczeretal01}. Other features are the spectral breaks at $\nup=\gammap^2\nu_0$, $\nucool=\gcool^2\nu_0$, and the exponential cut-off at $\numax=\gammamax^2\nu_0$. This model, therefore, allows a simple homogeneous source to reproduce the common features shown by many IDV sources. \begin{acknowledgement} We thank Luisa Ostorero and Stefan Wagner for helpful discussions and for providing us with easy access to the observational data. We would also like to thank the anonymous referee for constructive comments and suggestions that we feel have led to a significant improvement in this paper. \end{acknowledgement} | 7 | 10 | 0710.1564 |
0710 | 0710.1939_arXiv.txt | We present a new code aimed at the simulation of diffusive shock acceleration (DSA), and discuss various test cases which demonstrate its ability to study DSA in its full time-dependent and non-linear developments. We present the numerical methods implemented, coupling the hydrodynamical evolution of a parallel shock (in one space dimension) and the kinetic transport of the cosmic-rays (CR) distribution function (in one momentum dimension), as first done by Falle. Following Kang and Jones and collaborators, we show how the \emph{adaptive mesh refinement} technique (AMR) greatly helps accommodating the extremely demanding numerical resolution requirements of realistic (Bohm-like) CR diffusion coefficients. We also present the \emph{parallelization} of the code, which allows us to run many successive shocks at the cost of a single shock, and thus to present the first direct numerical simulations of linear and non-linear \emph{multiple} DSA, a mechanism of interest in various astrophysical environments such as superbubbles, galaxy clusters and early cosmological flows. | Diffusive Shock Acceleration (DSA) at supernova remnant blast waves is the favoured production mechanism for the production of the galactic cosmic-rays (CR). This theory, developed since the late 70s (see \citealt{Drury1983a} for a review), has now both strong theoretical and observational supports. The theoretical grounds of the model lie in the early ideas of Fermi (\citeyear{Fermi1949a}, \citeyear{Fermi1954a}): the regular Fermi acceleration mechanism (known as Fermi~I, the stochastic one being known as Fermi~II) can naturally explain the formation of a power-law spectrum by a shock wave -- with a remarkable universal slope whose value $s$ depends solely on the shock compression ratio $r$ (which is always 4 for strong non-relativistic shocks). However, the acceleration process can easily be so efficient that the CR may back-react on the shock dynamics, modifying the acceleration process in a fully non-linear way, and requiring a much more detailed analysis (see \citealt{Malkov2001c} for a review). Thus the DSA mechanism has still received a lot of attention in the last 20 years, from both a theoretical and a numerical perspective. Analytical works have been mostly limited to the test-particle (linear) regime. The full non-linear time-dependent problem has been mostly investigated through numerical simulations, using several different approaches (see \citealt{Jones2001a} for a short review). A first class is based on particle methods, from the early Monte-Carlo simulations developed by \cite{Ellison1984a} to the recent Particle-In-Cells codes (eg \citealt{Dieckmann2000a}). An alternate approach consists of solving the (Fokker-Planck) transport equation. This has first been done in the "two fluid" model (eg \citealt{Jones1990a}), then dealing with the complete particle distribution function (\citealt{Falle1987a}, \citealt{Bell1987a}, \citealt{Kang1991a}, \citealt{Duffy1992a}). Most of this work has been aimed at understanding the role of single (isolated) supernovae remnants, although in many contexts CR are likely to experience many shocks, most notably inside superbubbles (\citealt{Parizot2004a}). In this paper we present a new code for the study of DSA, named \emph{Marcos} for \emph{Machine {\`a} Acc{\'e}l{\'e}rer les Rayons COSmiques}\footnote{the French for \emph{COSmic-Rays Acceleration Machine}}. In section~\ref{sec-dsa} we present the basics of the numerical methods implemented in our code, which couples the hydrodynamical evolution of a fluid with the kinetic transport of the CR. In section~\ref{sec-scales} we present the Adaptive Mesh Refinement (AMR) technique which allows us to resolve the (very) different scales induced by CR diffusion. In section~\ref{sec-multi} we present parallelization of the code, to be able to study in reasonable wall-clock time the effects of multiple shocks. | We have presented a new code aimed at the simulation of time-dependent non-linear diffusive shock acceleration. It is based on the kinetic approach, coupling the hydrodynamical evolution of the plasma with the diffusive transport of the distribution function of the supra-thermal particles. As such it falls under the legacy of the pioneers (\citealt{Falle1987a}, \citealt{Duffy1992a}) and of the masters (\citealt{Kang1991a}, \citealt{Kang2001a}) of the genre. As the CRASH code it implements an efficient AMR technique to deal with the huge range of space- and time-scales induced by CR diffusion of Bohm-like type. To save even more on computing time we have also parallelized our code in momentum so that we can study acceleration by multiple shocks as fast as acceleration by a single shock. However in many aspects (high-Mach flows, shock tracking, self-consistent injection) our code remains numerically simpler than CRASH -- which can be both a limitation and an advantage. Regarding the physics we note that various mechanisms of importance could be included in the code: self-consistant diffusion coefficient (adding magnetic waves transport), second-order Fermi acceleration (especially between multiple shocks), electrons acceleration (adding radiative losses), CR radiation (hadronic and leptonic)\ldots We have presented a few tests that show that our code works well, with respect to both the physical accuracy and the numerical efficiency, even in realistic difficult situations. We are now able to investigate in details the various aspects of the DSA mechanism, which 30 years after its early developments still poses some difficulties. In particular we can address the non-linear \emph{multiple} DSA mechanism, which we believe hasn't received so far all the attention it deserves. Our very first results suggest that the injection fraction plays a crucial role. We intend now to study in more details the situation in superbubbles. % | 7 | 10 | 0710.1939 |
0710 | 0710.1108_arXiv.txt | We study the limits of accuracy for weak lensing maps of dark matter using diffuse 21-cm radiation from the pre-reionization epoch using simulations. We improve on previous ``optimal'' quadratic lensing estimators by using shear and convergence instead of deflection angles. We find that non-Gaussianity provides a limit to the accuracy of weak lensing reconstruction, even if instrumental noise is reduced to zero. The best reconstruction result is equivalent to Gaussian sources with effective independent cell of side length $2.0h^{-1}\, \rm Mpc$. Using a source full map from z=10-20, this limiting sensitivity allows mapping of dark matter at a Signal-to-Noise ratio (S/N) greater than 1 out to $l\lesssim 6000$, which is better than any other proposed technique for large area weak lensing mapping. | \label{INTRO} The lens mapping of dark matter is an essential cornerstone of modern precision cosmology. Weak gravitational lensing has developed rapidly over the past years, which allows the measurement of the projected dark matter density along arbitrary lines-of-sight using galaxies as sources. Recently, \citet{2007PhRvD..76d3510S} have demonstrated the first CMB lensing detection. The goal is now to achieve high precision cosmological measurements through lensing, at better than 1\% accuracy. Galaxies are plentiful on the sky, but their intrinsic properties are not understood from first principles, and must be measured from the data. Future surveys may map as many as $10^{10}$ source objects. Using galaxies as lensing sources has several potential limits \citep{2004PhRvD..70f3526H}, including the need to calibrate redshift space distributions and PSF corrections, to be better than the desired accuracy, say 1\%. This will be challenging for the next generation of experiments. Some sources, such as the CMB, are in principle very clean, since its redshift and statistical properties are well understood. Unfortunately, there is only one 2-D CMB sky with an exponential damping at $l \gg 1000$, which limits the number of source modes to $\sim 10^6$. The potential of detecting the 21-cm background from the dark ages will open a new window for cosmological detections. Studying the 21-cm background as high redshifts lensing source, as well as the physics of the 21-cm background itself, provide rich and valuable information to the evolution of universe. The number of modes on the sky is potentially very large, with numbers of $10^{16}$ or more. For this reason, 21-cm lensing has recently attracted attention. However, most of the reconstruction methods are based on a Gaussian assumption \citep{2004NewA....9..417P,2004NewA....9..173C,2006ApJ...653..922Z, 2006astro.ph.11862B,2007arXiv0706.0849H}. In contrast to CMB lensing, where the Gaussian assumption works well, non-Gaussianity in 21-cm lensing may affect the results. Non-linear gravitational clustering leads to non-Gaussianity, and ultimately to reionization. In this paper, we will address the problem of the lensing of pre-reionization gas. 21-cm emission is similar to CMB: both are diffuse backgrounds. It is natural to apply the techniques used in CMB lensing. \citet{2002ApJ...574..566H} expand the CMB lensing field in terms of the gravitational potential (or deflection angles), and construct a trispectrum based quadratic estimator of potential with maximum S/N. However, unlike CMB, the 21-cm background has a 3-D distribution and is intrinsically non-Gaussian. A fully 3-D analysis is explored in \citet{2006ApJ...653..922Z}, where they generalize the 2-D quadratic estimator of CMB lensing \citep{2002ApJ...574..566H} to the 3-D Optimal Quadratic Deflection Estimator (OQDE). A local estimator was proposed in \citet{2004NewA....9..417P}, which assumed a power law density power spectrum. In this paper, we will design localized estimators for the lensing fields under the Gaussian assumption, and apply the derived reconstruction technique to Gaussian and non-Gaussian sources. The influence of non-Gaussianity can be measured by comparing the numerical results between the Gaussian sources and non-Gaussian sources. Quadratic lensing reconstruction is a two point function of the lensed brightness temperature field of the 21-cm emission. In the paper, 3-D quadratic estimators are constructed for the convergence ($\kappa$), as well as the shear ($\gamma$). Our method recovers the $\kappa$ and $\gamma$ directly instead of gravitational potential or deflection angles. Our estimators have in principle the same form as the OQDE, consisting of the covariance of two filtered temperature maps. The OQDE reconstructs the deflection angle, while our estimators reconstruct the kappa and shear fields. Our filtering process can be written as a convolution of the observed fields. As presented in Appendix and section 4, our combined estimator is unbiased, and equally optimal as the OQDE for Gaussian sources, and has better performance for non-Gaussian sources, and recovers three extra (constant) modes. Other authors also developed reconstruction methods from alternative approaches. \citet{2006astro.ph.11862B} give a estimator for shear. They choose the separate 2-D slices at certain redshift intervals, and then these slices can be treated as independent samples for the same lensing structure. As a result, the information between these slices are lost. \citet{2004NewA....9..173C} expands the lensed field to a higher order of the gravitational potential, and investigates the higher order correction to the lensed power spectrum. The paper is organized as follows: The basic framework of lensing and the reconstruction method is introduced in $\S 2$. The numerical methods are presented in $\S 3$. The results are discussed in $\S 4$. We conclude in $\S 5$. | \label{CONC} In this paper, we developed the maximum likelihood estimator for the large-scale structure from the 21-cm emission of the neutral gas before the epoch of re-ionization. The convergence and shears can be constructed independently. To test the effects of non-Gaussianity, we applied our estimators to simulated data. The sources were generated by N-body simulations, because gas is expected to trace the total mass distribution. To investigate the influence of non-Gaussianity, we also use Gaussian sources which have the same power spectrum as the simulated sources. We applied our estimator and the OQDE on both the Gaussian and non-Gaussian sources. Though our estimators are derived in the simplified case of a constant convergence, the noise of our combined estimator of convergence and shear are the same as the OQDE for Gaussian sources. For a finite survey area, three extra constant modes can be recovered. The non-Gaussian nature of the source can increase the error bar by orders of magnitude, depending on the experimental cut off scale. Shear construction is affected less by non-Gaussianity than the convergence field, and the combined estimator with non-Gaussian noise weights is a better choice than reconstructing with the OQDE. S/N can not be boosted infinitely by reducing the experimental noise, and achieves its maximum for a cut off around $k^{\rm NG}_{\rm c}\approx 4h\,\Mpc^{-1}$. Below that scale the S/N start to saturate or even decrease. The maximum S/N for non-Gaussian sources is equal to Gaussian sources with $k^{\rm G}_{\rm c}\approx2h\,\Mpc^{-1}$, where the power spectrum of source is $\Delta^2\approx0.2$ and the side length of the effectively independent cells is $ 2.0 h^{-1}\,\rm Mpc$. The maximum S/N is greater than unity for $l\lesssim 6000$, which makes 21-cm lensing very competitive compared to optical approaches. {\it Acknowledgments} We thank Oliver Zahn, Chris Hirata, Brice M\'{e}nard and Mike Kesden for helpful discussions. T.T. Lu thanks Pengjie Zhang, Zhiqi Huang, Hy Trac, and Hugh Merz for help in the early stage of the work. | 7 | 10 | 0710.1108 |
0710 | 0710.0562_arXiv.txt | Several studies have correlated observations of impulsive solar activity --- flares and coronal mass ejections (CMEs) --- with the amount of magnetic flux near strong-field polarity inversion lines (PILs) in active regions' photospheric magnetic fields, as measured in line-of-sight (LOS) magnetograms. Practically, this empirical correlation holds promise as a space weather forecasting tool. Scientifically, however, the mechanisms that generate strong gradients in photospheric magnetic fields remain unknown. Hypotheses include: the (1) emergence of highly twisted or kinked flux ropes, which possess strong, opposite-polarity fields in close proximity; (2) emergence of new flux in close proximity to old flux; and (3) flux cancellation driven by photospheric flows acting fields that have already emerged. If such concentrations of flux near strong gradients are formed by emergence, then increases in unsigned flux near strong gradients should be correlated with increases in total unsigned magnetic flux --- a signature of emergence. Here, we analyze time series of MDI line-of-sight (LOS) magnetograms from several dozen active regions, and conclude that increases in unsigned flux near strong gradients tend to occur during emergence, though strong gradients can arise without flux emergence. We acknowledge support from NSF-ATM 04-51438. | \label{sec:intro} It has been known for decades that flares and filament eruptions (which form CMEs) originate along polarity inversion lines (PILs) of the radial photospheric magnetic field. In studies using photospheric vector magnetograms, Falconer {\em et al.}~(2003, 2006) \nocite{BW_Falconer2003,BW_Falconer2006} reported a strong correlation between active region CME productivity and the total length of PILs with strong potential transverse fields ($>150$ G) and strong gradients in the LOS field (greater than 50 G Mm$^{-1}$). They used a $\pm$2-day temporal window for correlating magnetogram properties with CMEs. Falconer {\em et al.}~(2003) \nocite{BW_Falconer2003} noted that these correlations remained essentially unchanged for ``strong gradient'' thresholds from 25 to 100 G Mm$^{-1}$. Using more than 2500 MDI (LOS) magnetograms, Schrijver (2007) \nocite{BW_Schrijver2007} found a strong correlation between major (X- and M-class) flares and the total unsigned magnetic flux near (within $\sim 15$ Mm) strong-field PILs --- defined, in his work, as regions where oppositely signed LOS fields that exceed 150 G lie closer to each other than the instrument's $\sim$ 2.9 Mm resolution. Schrijver's (2007) effective gradient threshold, {$\sim$ 100 G Mm$^{-1}$}, is stronger than that used by Falconer {\em et al.}~(2003, 2006). \nocite{BW_Falconer2003,BW_Falconer2006} Although these studies were published recently, the association between flares and $\delta$ sunspots, which posses opposite-sign umbrae within the same penumbra --- and therefore also possess strong-field PILs --- has been well known for some time \cite{BW_Kunzel1960,BW_Sammis2000}. In particular, $\beta \gamma \delta$ spot groups are most likely to flare \cite{BW_Sammis2000}. A $\beta \gamma$ designation means no obvious north-south PIL is present in an active region \cite{BW_Zirin1988}. We note that Cui {\em et al.}~(2006) \nocite{BW_Cui2006} found that the occurrence of flares is correlated with the maximum magnitude of the horizontal gradient in active region LOS magnetograms --- not just near PILs --- and that the correlation increases strongly for gradients stronger than $\sim$ 400 G Mm$^{-1}$. One would expect the measures of CME- and flare- productivity developed by both Falconer {\em et al.}~(2003,2006) \nocite{BW_Falconer2003,BW_Falconer2006} Schrijver (2007) \nocite{BW_Schrijver2007} to be larger for larger active regions. Importantly, however, both studies showed that their measures of flux near strong-field PILs is a better predictor of flare productivity than total unsigned magnetic flux. Evidently, more flux is not, by itself, as significant a predictor of flares as more flux near strong-field PILs. These intriguing results naturally raise the question, ``How do strong-field PILs form?'' For brevity, we hereafter refer to strong-field PILs as SPILs. Schrijver (2007) \nocite{BW_Schrijver2007} contends that large SPILs form primarily, if not solely, by emergence. But he also noted that flux emergence, by itself, does not necessarily lead to the formation of SPILs. Rather, a particular type of magnetic structure must emerge, one containing a long SPIL at its emergence. He suggests such structures are horizontally oriented, filamentary currents. Beyond the ``intact emergence'' scenario presented by Schrijver (2007), \nocite{BW_Schrijver2007} other mechanisms can generate SPILs. When new flux emerges in close proximity to old flux --- a common occurrence \cite{BW_Harvey1993} --- SPILs can form along the boundaries between old and new flux systems. Converging motions in flux that has already emerged can also generate SPILs. If the convergence leads to flux cancellation by some mechanism --- emergence of U loops, submergence of inverse-U loops, or reconnective cancellation \cite{BW_Welsch2006} --- then the total unsigned flux in the neighborhood of the SPIL might decrease as the SPIL forms. We note that, while cancellation in already-emerged fields can occur via flux emergence (from upward moving U-loops), the emergence of a new flux system across the photosphere must increase the total unsigned flux that threads the photosphere. If the emergence of new flux were primarily responsible for SPILs, then a straightforward prediction would be that an increase in total unsigned flux should be correlated with an increase in the amount of unsigned flux near SPILs. Hence, observations showing that increases in the unsigned flux near SPILs frequently occur without a corresponding increase in total unsigned flux would rule out new flux emergence as the sole cause of these strong field gradients. Our goal is to investigate the relationship between increases in the amount of unsigned flux near SPILs with changes in unsigned flux in the active regions containing the SPILs, to determine, if possible, which processes generate SPILs. | In Figure \ref{fig:tres}, we show a scatter plot of changes in $R$ as a function of changes in ${\cal B}$. The plot does not show the full range in $\Delta {\cal B}$, but the $\Delta R$ for outliers on the horizontal axes are near zero. One striking feature of the plot is its flatness, i.e., that most changes in ${\cal B}$ are not associated with any change in $R$. In Table \ref{tab:uno}, we tabulated the data points in each quadrant of this plot. Clearly, increases in $R$, the unsigned flux near SPILs, usually occur simultaneously with increases in the unsigned flux over the entire active region. Increases in $R$ only occur less frequently when flux is decreasing, i.e., during cancellation. \begin{figure}[!ht] \includegraphics[width=5.5in]{welb_fig03.eps} \caption{A scatter plot of changes in $R$ as a function of changes in ${\cal B}$. Increases in $R$, the unsigned flux near SPILs, usually occur simultaneously with increases in the unsigned flux over the entire active region. Increases in $R$ only occur rarely when flux is decreasing, i.e., during cancellation. For a breakdown of the data points in each quadrant, see Table \ref{tab:uno}.} \label{fig:tres} \end{figure} \begin{table} \caption{Breakdown of Flux Changes \label{tab:uno}} \begin{tabular}{l|c|c} & $\Delta {\cal B} < 0$ & $\Delta {\cal B} > 0$ \\ \hline $\Delta R > 0$ & 215 & 671 \\ \hline $\Delta R < 0$ & 363 & 371 \\ \end{tabular} \end{table} We set out to answer the question, ``How do strong-field PILs form?'' We related changes in total, unsigned flux over whole active regions with changes in total, unsigned flux in subwindows of the same active regions --- defined by weighting maps. One might expect, therefore, that these quantities should be correlated, casting doubt about our ability to discrimintate between changes in total flux in active regions and in subwindows. If the two were strongly correlated, the excess of events with $\Delta R > 0$ and $\Delta {\cal B} > 0$ might not be very meaningful. In fact, however, $\Delta R$ and $\Delta {\cal B}$ are poorly correlated: the two have a linear correlation coefficient $r = 0.29$, and a rank-order coefficient of $0.36.$ This suggests that the relationship between increases in $R$ and increases in total, unsigned active region flux is not an artifact of our approach. Nonetheless, our active region sample is not ideally suited to address the origin of SPILs, generally. Our sample was not unbiased with respect to active region morphology; we selected regions with well-defined PILs. In addition, our sample included some decayed active regions that NOAA AR designations. Consequently, we believe that a follow-up study, with a much larger, unbiased sample of active regions, is warranted. With caveats, therefore, our study supports Schrijver's (2007) \nocite{BW_Schrijver2007} contention that the emergence of new flux creates the strong-field polarity inversion lines that he found to be correlated with flares. | 7 | 10 | 0710.0562 |
0710 | 0710.5171_arXiv.txt | We describe a method for computing the biases that systematic signals introduce in parameter estimation using a simple extension of the Fisher matrix formalism. This allows us to calculate the offset of the best fit parameters relative to the fiducial model, in addition to the usual statistical error ellipse. As an application, we study the impact that residual systematics in tomographic weak lensing measurements. In particular we explore three different types of shape measurement systematics: (i) additive systematic with no redshift evolution; (ii) additive systematic with redshift evolution; and (iii) multiplicative systematic. In each case, we consider a wide range of scale dependence and redshift evolution of the systematics signal. For a future DUNE-like full sky survey, we find that, for cases with mild redshift evolution, the variance of the additive systematic signal should be kept below $10^{-7}$ to ensure biases on cosmological parameters that are sub-dominant to the statistical errors. For the multiplicative systematics, which depends on the lensing signal, we find the multiplicative calibration $m_0$ needs to be controlled to an accuracy better than $10^{-3}$. We find that the impact of systematics can be underestimated if their assumes redshift dependence is too simplistic. We provide simple scaling relations to extend these requirements to any survey geometry and discuss the impact of our results for current and future weak lensing surveys. | \label{intro} Weak gravitational lensing, or `cosmic shear', is undergoing a phase of rapid expansion \citep[see][for reviews]{2003ARA&A..41..645R,2006astro.ph.12667M,2003astro.ph.10908H} with many future surveys and instruments being planned (e.g. DUNE\footnote{http://www.dune-mission.net}, PanSTARRS\footnote{http://pan-starrs.ifa.hawaii.edu}, DES\footnote{https://www.darkenergysurvey.org}, SNAP\footnote{http://snap.lbl.gov} and LSST\footnote{http://www.lsst.org}). Central to the planning and designing of these instruments is our ability to predict the uncertainties that such measurements will achieve on the cosmological parameters. To this end, the Fisher matrix has become a widely used tool in cosmology for calculating their covariance matrix. However, a limitation of this approach is that it is only able to account for statistical errors, i.e. ones that cause an enlargement of the error bars, and is not well-suited for treatment of systematic errors, which can introduce biases that move the measured central value relative to its true value. One approach that is commonly taken to overcome this limitation is to treat the systematic errors in the same way as statistical errors and to marginalise over possible values. This introduction of nuisance parameters, in addition to the cosmological parameters, causes the error ellipses to expand. A more accurate approach, which we use here, is to directly calculate the bias that the systematic signals will introduce. This bias will tend to offset the central value to the measurements from the true values, as shown in figure \ref{fig:fig0}. Since the computations needed for this calculation are very similar to those performed in the standard Fisher matrix analysis, extending the current Fisher matrix analysis to include a calculation of bias is relatively straightforward. We apply this formalism to study the impact of residual systematics on tomographic cosmic shear surveys. In particular, we consider systematics arising in the measurement of galaxy shapes after correction of instrumental effects (such as the Point Spread Function). We consider both additive and multiplicative systematics and explore a wide range of scale and redshift dependences. In an earlier work, \cite{2006MNRAS.366..101H} considered the impact of photometric calibration errors and power law shape systematics using the Fisher matrix formalism. A similar bias formalism was introduced by \cite{2005APh....23..369H} and applied to theoretical uncertainties in modeling the matter power spectrum with N-body simulations. Our work expands upon these earlier works, by appying the bias formalism to a broad set of shape systematics and by studying the joint impact of systematic and statistical errors in current and future surveys. This paper is organised as follows. In section 2, we describe the formalism that we use to quantify systematic biases. In section 3, we apply our formalism to cosmic shear surveys by exploring the effect of three types of shape measurement systematics: (i) additive with no redshift evolution; (ii) additive with redshift evolution; and (iii) multiplicative. For each type, we consider several possibilities for their scale dependence: (i) log-linear systematics; (ii) systematics that have the same shape as the lensing signal; and (iii) systematics that mimic a small change in the cosmological parameters. In section 4, we study the impact of the systematics in the design of future surveys. Our conclusions are summarised in section 5. | \label{conclusion} In this paper, we have outlined a method for computing the biases that residual systematics introduce. This approach involves a simple extension of the Fisher matrix formalism that is now widely used in cosmology to make error forecasts. As an application, we have used it to study the impact that residual systematic signals will have on future tomographic cosmic shear measurements. Specifically, we have explored three different types of shape systematic signal affecting tomographic shear power spectra: (i) additive systematics with no redshift evolution; (ii) additive systematics with redshift evolution; and (iii) multiplicative systematics. The requirement target is then to have all types of systematics close to zero. This defines a tolerance envelope for the systematics that allows the residual systematics errors in the power spectra to be positive or negative within its limits. It is important to note that it is the $worst$ systematic possible within this limit which drives the requirement, not a marginalised average over all systematics. To this end we have investigated a wide class of possible systematic shapes and used the most constraining ones to set our systematic requirements. In doing this, we have found that, for both the additive and multiplicative parts, it is vital to consider systematics that have positive and negative power spectra $C_\ell^{sys}$. For instance we see, in the multiplicative case, that investigating only power-law behaviour for its redshift evolution (i.e. $m$ is always positive) can lead to a factor of 5 underestimation of the impact of a systematic within a given tolerance window. From our calculation we are able to set the following requirements on the survey we have considered (a DUNE-like survey covering 20,000 sq. degrees with 35 galaxies per arcmin$^{-2}$ and a median redshift of $z_m=0.9$): \begin{itemize} \item For both the additive and the multiplicative signals, the redshift evolution needs to be weak $\beta < 1.5$, where the errors for a given galaxy scale as $(1+z)^\beta$. This is not a trivial requirement since the shapes of more distant galaxies are harder to measure since they are smaller and fainter. \item The power spectrum of the residual additive shear error, that is the part that is not correlated to the lensing signal, must be controlled such that its amplitude is $\sigma^2_{sys} < 10^{-7}$ (as defined in equation \ref{eq:var}) \item The multiplicative part needs to be controlled to a precision of $m_0 <10^{-3}$, where $m_0$ is the shear calibration error. This means that we need to be able calibrate shears to an accuracy of 0.1\%, which is about one order of magnitude better than the current best measurement methods are able to achieve, as determined by the latest STEP simulations \citep{2006MNRAS.368.1323H,2007MNRAS.376...13M}. \end{itemize} These specific requirements apply to our fiducial survey, but we provide scaling relations (Eqs \ref{eq:scale1} and \ref{eq:scale2}) which show the requirements for any survey geometry. We have shown that for current survey covering $\sim 100$ deg$^{2}$ \citep{2007MNRAS.381..702B}, we need $\sigma^2_{sys} < 3\times10^{-6}$ and $m< 0.03$. This level of accuracy is at the limit of the performance of the current best shear measurement methods, as demonstrated by STEP. However, further systematics, such as that arising from PSF calibration and interpolation, are not accounted for by STEP and can dominate the error budget for future surveys. A discussion of the requirements for additive systematics and PSF modeling in the context of present and future surveys will be presented in a later paper (Paulin-Henriksson et al., 2007, in prep). \newpage \appendix | 7 | 10 | 0710.5171 |
0710 | 0710.3754.txt | We present arcsecond-scale mid-ir photometry (in the 10.5 $\mu$m N band and at 24.8 $\mu$m), and low resolution spectra in the N band ($R\simeq100$) of a candidate high mass protostellar object (HMPO) in IRAS 18151-1208 and of two HMPO candidates in IRAS 20343+4129, IRS 1 and IRS 3. In addition we present high resolution mid-ir spectra ($R\simeq80000$) of the two HMPO candidates in IRAS 20343+4129. These data are fitted with simple models to estimate the masses of gas and dust associated with the mid-ir emitting clumps, the column densities of overlying absorbing dust and gas, the luminosities of the HMPO candidates, and the likely spectral type of the HMPO candidate for which [Ne II] 12.8 $\mu$m\ emission was detected (IRAS 20343+4129 IRS 3). We suggest that IRAS 18151-1208 is a pre-ultracompact HII region HMPO, IRAS 20343+4129 IRS 1 is an embedded young stellar object with the luminosity of a B3 star, and IRAS 20343+4129 IRS 3 is a B2 ZAMS star that has formed an ultracompact HII region and disrupted its natal envelope. | Many open questions in high-mass star formation are related to the evolution of circumstellar envelopes, accretion disks, and jets from high-mass protostellar objects (HMPOs). HMPOs are often bright sources in mid-ir continuum, but only in a few recent cases have images suggested specific structures such as disks, jets, or warm outflow cavity walls \citep{sri05, deb05, deb06,deb07}. Mid-ir ionic lines like [Ne II] and [S IV] have been used to map compact and ultracompact HII regions (UC HII) and photodissociation regions, and to study their structure and excitation \citep{lac82,oka01,zhu05,kas02,kas06}, but there is a lack of observations of HMPOs. A remarkable feature of high-mass star formation is that HMPOs, defined as actively accreting mass, can begin nuclear fusion and hence also be rapidly evolving massive young stellar objects (MYSOs) that have already formed hypercompact or ultracompact HII regions \citep{beu07}. This feature raises a possibility of determining the spectral type of an MYSO through the ionic lines excitation, or from the number of ionizing photons required for its observed centimeter continuum flux, separately from estimating its luminosity and spectral type from infrared emission. However there is also the possibility that the ionization is collisionally excited by a jet. It may be possible to distinguish between the two cases, depending on the Doppler velocities, the morphology of ionized gas, and the ratio of the required flux of ionizing photons to total luminosity. Hoping to enlarge the sample of HMPOs that could be studied in detail despite potential limitations, in 2003 we made mid-ir observations on the IRTF of about a third of the survey of 69 HMPO candidates presented by \citet{sri02} and \citet{beu02a}. We chose objects from their survey that appeared compact and/or bright in the MSX survey, and found that about 80\% of them were unresolved or marginally resolved by MIRSI \citep{deu02} on the IRTF in the broad N band at 10.5 $\mu$m and in a narrow band filter at 24.8 $\mu$m. In addition, we obtained MIRSI grism low-resolution spectra (R $\simeq100$) of ten of them in the N band. In 2006 on Gemini North, we obtained TEXES \citep{lac02} high-resolution spectra (R $\simeq80000$) of two HMPO candidates for which we had grism spectra. In this paper we present spectra and photometry of three candidate HMPOs including the two with TEXES spectra: IRAS 18151-1208, IRAS 20343+4129 IRS 1, and IRAS 20343+4129 IRS 3 (hereafter, 18151, 20343 IRS 1, and 20343 IRS 3). We will demonstrate that mid-ir emission from the dust and gas near the HMPO candidates (where it is strongly heated) can be used as a useful probe of temperatures, masses, and luminosities, using simple isothermal clump models, even if each component (envelope, disk, jet, or cavity wall) is not resolved. In combination with observations cited below, we are able to use our new data to infer the nature of each candidate candidate HMPO (e.g. pre-UCHII region HMPO, ZAMS B2 star). The objects chosen are near the centers of complex, large-scale massive molecular outflows mapped by \citet{beu02b}. 20343 has an apparent large-scale N-S outflow whose red and blue lobes are both extended E-W \citep{beu02b}, but IRS 1 also has a compact E-W velocity outflow in CO(2-1), while IRS 3 presents an ambiguous situation \citep{pal06}. All objects show near-ir emission from shocked H$_2$ \citep{dav04,kum02}. Two of them (18151 and 20343 IRS 3) were observed to have 0.5 and 1.8 mJy 3.6 cm emission, respectively \citep{car99,sri02}. We observed 20343 IRS 1 and IRS 3 with TEXES on Gemini North based on the 3.6 cm and H$_2$ observations, with the goal of studying the role of ionized gas in them. The 10 $\mu$m grism spectral shapes of the HMPO candidates fall into three classes: those with deep silicate absorption; those with moderate silicate absorption and an apparent peak at about 8.5 $\mu$m; and those without an apparent silicate absorption feature but with continuum rising monotonically from short to long wavelengths (Campbell et al. 2007, in preparation). Examples of these shapes can be seen in the UC HII spectra presented by \citet{fai98}. Since the HMPO candidates were chosen based on IRAS colors similar to UC HII regions \citep{sri02}, one would expect the HMPO candidates to have similar 10 $\mu$m spectra. IRAS 20343+4129, was observed with the IRAS LRS and has a silicate absorption feature \citep{vol91}. The three objects presented here include an example of each of the three classes of grism spectral shapes. Two of the ten candidate HMPO grism spectra show strong [Ne II] lines, IRAS 18247-1147 and 18530+0215. \citet{sri02} reported relatively strong 3.6 cm fluxes of 47 and 311 mJy, respectively, for them. We did not observe them with TEXES on Gemini in order to focus on the earliest possible HMPO stage associated with ionized gas. Deriving physical information from the continuum spectra is difficult because the actual geometry of the dust distribution is unknown, except that the sizes of the N band and 24.8 $\mu$m images limit the extent of the emitting structures. Experience with one-dimensional radiative transfer models of spectral energy distributions from UC HII regions \citep{cam95,cam00,cam04}, and inspection of the new spectra themselves suggest that there are ranges of temperatures in the emitting regions. However, the two-dimensional radiative transfer models of \citet{deb05b} and \citet{whi03a,whi03b,whi04} show that orientation of flattened envelopes and outflow cavities dramatically affects the depth of the silicate absorption feature, as does emission and absorption by individual clumps in a clumpy dust cloud in the three-dimensional models of \citet{ind06}. The recent observations of extended and complex near- and mid-ir emission from HMPOs \citep{deb05,deb06,sri05} also indicate that one-dimensional models are unrealistic. Nevertheless, a simple three part model will allow us to derive approximate temperatures, column densities, and masses of different dust components. The first component is hot dust that could be (part of) a relatively compact disk, or a clump very near the HMPO candidate; the second is warm dust that could be a more extended (part of a) disk, a clump of dust further out, or perhaps the inner wall of of an outflow cavity; and the third is cold dust in an outer envelope that creates the silicate absorption feature. Deriving information from the ionic lines of an UC HII region is in principle straight-forward. The line fluxes can be corrected for local extinction based on the continuum models discussed above, and then the ratio of [Ne II] to [S IV] fluxes can be used to determine the exciting star's temperature \citep{lac82,oka01}. The numerical simulation code cloudy (http://www.nublado.org/) can also be used to deduce the star's temperature and luminosity by matching the simulated intensity and spatial extent of free-free, [Ne II], and [S IV] emission to the observations. The star's temperature can then be compared to that of the spectral type deduced from the cm continuum flux. In addition, comparison of Doppler velocities of the ionic lines to those of molecular lines can be used to indicate if the gas is in an UC HII region or a jet. | We have presented high resolution mid-ir observations made with MIRSI on the IRTF and TEXES on Gemini North of three HMPO candidates taken from a partial follow up survey of HMPO candidates originally studied at 1.2 mm and radio wavelengths by \citet{sri02} and \citet{beu02a}. They are typical for HMPO candidates observed in the follow up survey being compact at 1$\arcsec$ resolution, having low resolution spectra with strong, moderate, or weak silicate absorption, and with one emitting the [Ne II] line. A simple model of hot dust in emission, warm dust in emission, and cold dust in absorption was developed to fit our 8-13 $\mu$m low resolution spectra and our 24.8 $\mu$m photometric points. Even an apparently flat 8-13 $\mu$m spectrum requires an absorption component if the underlying emission is assumed to be due to hot or warm silicate dust. The temperatures ranged from $\sim$ 400-1000 K for the hot dust, and $\sim$ 100-200 K for the warm dust. Using \citet{dra03a} $R_V$=5.5 model dust properties and gas-to-dust ratio, only small masses of gas and dust in the two emitting components are needed to fit the data. The masses are less than about 1/10 solar mass (often much less) even though these are high or intermediate mass stars, and the mid-ir emission cannot be due the the bulk of the mass in massive accretion disks. The mid-ir is likely to be emitted by the inner walls of outflow cavities and perhaps partly by the surfaces of accretion disks. On the other hand, high column densities, $10^{22} - 10^{23}$ H nucleons cm$^{-2}$, are required for the cold absorption components. These column densities are less than derived from 1.2 mm 11$\arcsec$ data using \citet{dra03a} dust, but the discrepancy may be resolved if the slope of the absorption coefficient from far-ir to mm is flattened, as suggested by some observations. Our three component model is not meant to fit either near-ir or far-ir to mm ends of SEDs. Nevertheless, the dust we are modeling in the hot and warm components appears to absorb the bulk of an HMPO's or intermediate mass YSO's photospheric emission, so that the integrated flux of the two model components without application of the cold dust's extinction matches the luminosity as measured including the far-ir by IRAS. The mid-ir measurements together with the model thus give a reasonable way to determine the luminosity for individual HMPOs. The mid-ir emission of IRAS 18151-1208 together with weak 3.6 cm emission and other previous observations suggest that it is an early stage pre-UC HII HMPO whose luminosity is that of a B0 star. TEXES high resolution spectra that cover emission lines from ionized gas can be used to determine the nature of the emission (jet or HII region) and help determine the properties of the underlying star. In the case of IRAS 20343+4129 IRS 1, a lack of [NeII] emission, a well defined compact CO outflow, a moderately strong silicate absorption feature, and a dust model-based luminosity of 1400 L$_\sun$ imply that it is an intermediate mass YSO whose luminosity is that of a B3 star. For IRAS 20343+4129 IRS 3 observed [Ne II] emission and 3.6 cm free-free emission are consistent with a cloudy model indicating that the object is a B2 ZAMS star. Its weak silicate absorption and small mid-ir based luminosity suggest that it has already disrupted much of its natal envelope. % For IRAS 20343+4129 IRS 3, observed [Ne II] emission and % 3.6 cm free-free emission are consistent with a cloudy model % indicating that the object is a B2 ZAMS star. Its weak % silicate absorption and small mid-ir based luminosity suggest that it has already disrupted % much of its natal envelope. %% If you wish to include an acknowledgments section in your paper, %% separate it off from the body of the text using the | 7 | 10 | 0710.3754 |
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