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0710 | 0710.4301_arXiv.txt | We develop a formalism for studying the dynamics of massive black hole binaries embedded in gravitationally-bound stellar cusps, and study the binary orbital decay by three-body interactions, the impact of stellar slingshots on the density profile of the inner cusp, and the properties of the ejected hypervelocity stars (HVSs). We find that the scattering of bound stars shrinks the binary orbit and increases its eccentricity more effectively than that of unbound ambient stars. Binaries with initial eccentricities $e\gtrsim 0.3$ and/or unequal-mass companions ($M_2/M_1 \lesssim 0.1$) can decay by three-body interactions to the gravitational wave emission regime in less than a Hubble time. The stellar cusp is significantly eroded, and cores as shallow as $\rho\propto r^{-0.7}$ may develop from a pre-existing singular isothermal density profile. A population of HVSs is ejected in the host galaxy halo, with a total mass $\sim M_2$. We scale our results to the scattering of stars bound to Sgr A$^*$, the massive black hole in the Galactic Center, by an inspiraling companion of intermediate mass. Depending on binary mass ratio, eccentricity, and initial slope of the stellar cusp, a core of radius $\sim 0.1$ pc typically forms in 1-10 Myr. On this timescale about 500-2500 HVSs are expelled with speeds sufficiently large to escape the gravitational potential of the Milky Way. | In the standard paradigm of cosmic structure formation, it is expected that many wide massive black hole binaries (MBHBs) will form following the mergers of two massive galaxies (e.g. Begelman, Blandford, \& Rees 1980; Volonteri, Haardt, \& Madau 2003; Mayer \etal 2007). The binary will subsequently shrink due to stellar or gas dynamical processes and may ultimately coalesce by emitting a burst of gravitational waves. It was first proposed by Ebisuzaki, Makino, \& Okumura (1991) that the heating of the surrounding stars by a decaying SMBH pair would create a low-density core out of a preexisting cuspy (e.g. $\rho\propto r^{-2}$) stellar profile. In a purely stellar background a `hard' binary shrinks by capturing the stars that pass close to the holes and ejecting them at much higher velocities, a super-elastic scattering process (`gravitational slingshot') that depletes the nuclear region. Observationally, there is evidence in early--type galaxies for a systematically different distribution of surface brightness profiles, with faint ellipticals showing steep power-law profiles (cusps), while bright ellipticals have much shallower stellar cores (e.g. Faber \etal 1997; Ravindranath \etal 2001). Dwarf ellipticals seem to elude this paradigm showing somewhat flat profiles, similarly to bright ones (Graham \& Guzman 2003). Late--type spirals tend to show steep central cusps, as in the case of the Milky Way or Andromeda. Detailed N-body simulations have confirmed the cusp-disruption effect of a hardening MBHB (Makino \& Ebisuzaki 1996; Quinlan \& Hernquist 1997; Milosavljevic \& Merritt 2001), while semi-analytic modelling in the framework of hierarchical structure formation theories has shown that the cumulative damage done to a stellar cusps by decaying black hole pairs may explain the observed correlation between the `mass deficit' (the mass needed to bring a flat inner density profile to a $r^{-2}$ cusp) and the mass of the nuclear black hole (Volonteri, Madau, \& Haardt 2003). This is the third paper of a series aimed at the detailed study of the interaction between MBHBs and their stellar environment. In Sesana, Haardt, \& Madau (2006,2007a, hereafter Paper I and Paper II), we analyzed the three-body scatterings between a MBHB and background stars {\it unbound} to the binary. The assumption of a fixed background breaks down once the binary has ejected most of the stars on intersecting orbits, and the extraction of energy and angular momentum from the binary can continue only if new stars can diffuse into low-angular momentum orbits (refilling the binary's phase-space ``loss cone''), or via gas processes. In galaxies with inner cores or shallow cusps, only a small fraction of the loss cone is confined within the sphere of influence of the binary, and the approximation of a background of unbound stars is reasonable. A similar argument holds also in the case of a galaxy with a steep cusp hosting a nearly equal-mass binary ($M_1\sim M_2$). The radius of influence of such a pair, $r_{\rm inf}=G(M_1+M_2)/(2\sigma^2)$ where $\sigma$ is the stellar velocity dispersion, is of the order of the binary hardening radius, $a_h=GM_2/4\sigma^2$, and only few low-angular momentum stars have orbits with semi-major axis $\lesssim r_{\rm inf}$. This is not true for unequal-mass binaries, where $r_{\rm inf} \gg a_h$, and almost all interacting stars are bound to $M_1$. In this paper we develop a formalism for studying the dynamics of MBHBs embedded in gravitationally-bound stellar cusps. The plan is as follows. In \S~2 we describe our suite of three-body scattering experiments between the black hole pair and ambient bound stars. Our numerical results are used in \S~3 and \S~4 to construct a ``hybrid model'' of binary dynamics and investigate the orbital decay and shrinking of MBHBs in time-evolving stellar cusps. The properties of the ejected HVSs are discussed in \S~5. The massive black hole (Sgr A$^*$) in the Galactic Center and the stars around it offer a unique opportunity to study stellar dynamics in the extreme environment around a relativistic potential. The scattering of stars bound to Sgr A$^*$ by an inspiraling intermediate-mass black hole (IMBH) is treated in \S~6. Finally, we present a brief summary in \S~7. | We have performed, for the first time, scattering experiments between a MBHB and stars drawn from a cusp bound to the primary hole. We have studied the dynamics of the pair and its orbital decay by three-body interactions, the impact of the gravitational slingshot on the stellar density profile, the properties of the ejected stellar population, and have scaled our results to the case of Sgr A$^*$. Our results can be quickly summarized as follows: \begin{enumerate} \item The extraction of the cusp binding energy causes the binary to shrink by a larger factor compared to the scattering of unbound stars. The effect is more noticeable in the case of small mass ratios $q$. \item The binary orbital eccentricity increases much more rapidly compared to the unbound case. The eccentricity growth is more pronounced in small mass-ratio binaries, and for shallower stellar cusps. \item The combined effects of enhanced orbital decay and eccentricity growth lead very unequal-mass binaries to the gravitational wave coalescence phase. The detailed fate of the pair depends on the absolute value of its mass. More massive binaries decay faster. \item The stellar cusp is eroded, and the total mass removed by strong three--body encounters is 2-to-4 times the mass of the secondary hole. While the mass deficit caused by dynamical friction in a merger event scales with the binary mass (and involves mostly distant stars), the mass of stars ejected from the inner cusp by highly energetic interactions scales with $M_2$. Ejection occurs in a ``burst" lasting from few tenths to several thousands binary orbital periods, depending upon $q$. \item Scaled to the scattering of stars bound to Sgr A$^*$ by an inspiralling IMBH, our results imply the formation of a core of 0.1 pc in 1-10 Myrs, as well as the ejection of 500-2500 HVSs moving with speeds sufficient to escape the gravitational field of the Milky Way. In Sesana et al. (2007b) we have used the Brown et al. (2007) sample of unbound and bound HVSs together with numerical simulations of the propagation of HVSs in the Milky Way halo to constrain this ejection mechanisms, and shown that it appears to produce a spectrum of ejection velocities that is too flat compared to the observations. Future astrometric (as, e.g., {\it GAIA}) and deep wide-field (as, e.g., {\it LSST}) surveys should unambiguously identify the ejection mechanism of HVSs, and probe the Milky Way potential on scales as large as $200$ kpc (Gnedin et al. 2005; Yu \& Madau 2007). \end{enumerate} | 7 | 10 | 0710.4301 |
0710 | 0710.2900_arXiv.txt | We consider the possibility of cluster membership for 13 planetary nebulae that are located in close proximity to open clusters lying in their lines of sight. The short lifetimes and low sample size of intermediate-mass planetary nebulae with respect to nearby open clusters conspire to reduce the probability of observing a true association. Not surprisingly, line of sight coincidences almost certainly exist for 7 of the 13 cases considered. Additional studies are advocated, however, for 6 planetary nebula/open cluster coincidences in which a physical association is not excluded by the available evidence, namely M 1-80/Berkeley 57, NGC 2438/NGC 2437, NGC 2452/NGC 2453, VBRC 2 \& NGC 2899/IC 2488, and HeFa 1/NGC 6067. A number of additional potential associations between planetary nebulae and open clusters is tabulated for reference purposes. It is noteworthy that the strongest cases involve planetary nebulae lying in cluster coronae, a feature also found for short-period cluster Cepheids, which are themselves potential progenitors of planetary nebulae. | For some time our knowledge of the intrinsic properties of the Galaxy's population of individual planetary nebulae has been restricted by large uncertainties in their derived distances. \citet{zh95} suggests that the {\it average} uncertainty in the distances cited to Galactic planetary nebulae is in the range 35-50\%. Others are less optimistic. Such a large scatter may not be surprising, given that planetary nebulae exhibit various morphologies and span a large range in mass \citep{kw05}. In contrast, well-studied open clusters have distances and reddenings that are established to much greater precision, with distance uncertainties as small as 2.5\% being possible \citep{tb02}. Planetary nebulae established as members of open clusters are therefore a potential alternative means of calibrating their fundamental properties. With an inferred distance from cluster membership in conjunction with a planetary's angular diameter and expansion velocity, its true dimensions and age can be deduced. Cluster membership has the potential for a more direct calibration of the core mass-nebular He, C, and N abundance relationship expected in planetary nebulae as a result of single star evolution with asymptotic giant branch dredge-up \citep{ka00,cm00}. Planetary nebulae confirmed as cluster members would enhance their importance as calibrators for the Shklovsky relation \citep{of06} or other similar methods used to establish their distances \citep{bl01}. On a cautionary note, significant improvement in such relationships may not be possible if the observed scatter is intrinsic. Several factors conspire to reduce the probability of observing a planetary nebula associated with an open cluster. First, the effective sample of planetary nebulae includes a large number of objects that appear to populate the Galactic bulge \citep{as71,zi75}, according to catalogue statistics \citep{ko01} on their distribution along the Galactic plane (Figure 1), as well as their observed radial velocities \citep{of06}. Potential calibrators lying in nearby open clusters are greatly reduced in number when that population is excluded, although many spatial coincidences still exist \citep{zi75}. Associated open clusters with ages of less than $\sim28\times 10^6$ years ($\log(\tau)\leq7.5$) are likely to be excluded, since stellar evolutionary models indicate that the end products of their evolved components are Type II supernovae explosions. Current knowledge of stellar evolution suggests that the immediate precursors of C/O white dwarfs were planetary nebulae central stars that did not undergo core carbon ignition. In addition to a small sample size, the detection of an association between a planetary nebula and an open cluster is further hampered by the short lifetimes of planetary nebulae. Models indicate that their main-sequence progenitors were stars of $1-6.5\;M_{\sun}$ \citep[e.g.,][]{we00}, with an upper limit of $\sim8\;M_{\sun}$ being possible for production of Ne white dwarfs. The lifetime of the planetary nebula stage is very sensitive to initial progenitor mass and subsequent mass loss \citep[e.g.,][]{sb96}, and varies significantly for main-sequence turnoff ages greater than $\sim28\times 10^6$ years, with estimates ranging from $10^3$ to $10^5$ years \citep{sb96,ka00}. The most common age for nearby Galactic open clusters is $\sim100\times 10^6$ years ($\log(\tau) \simeq 8$), according to the catalogue compilation of \citet{di02} summarized in Figure 2. That corresponds to a main-sequence turnoff mass of $M_{TO}\simeq 4\;M_{\sun}$. The lifetime of planetary nebulae associated with such progenitors is of order $10^3$ years, essentially instantaneous on the Galactic stage. It is of interest to note that many planetary nebulae with massive central stars are found in the field, which is populated by the remnants of dissolved open clusters. Such clusters exceed the number of bound open clusters by a sizeable order \citep{la03}, which suggests that, despite the short lifetimes of planetary nebulae with massive central stars, increasing the statistical sampling of possible spatial coincidences between planetaries and clusters may lead to successful detections. The usefulness of such surveys at extragalactic scales by \citet{lr06} and \citet{ma06} is therefore obvious: larger statistics dominate and planetary nebulae are readily discernable, as demonstrated by their success as standard candles \citep{ja89}. The success of the Macquarie/AAO/Strasbourg H$\alpha$ (MASH) survey \citep{pa06} in detecting large numbers of additional Galactic planetary nebulae has also been extremely useful in revealing additional coincidences with Galactic clusters. The discovery of planetary nebulae within globular clusters \citep{ja97} raises a pertinent point that must be considered. If we consider $1-1.5\;M_{\sun}$ as a strict lower mass limit for the progenitors of planetary nebulae \citep{kw05,of06}, then, for the ages assigned to globular clusters, corresponding to main sequence turnoffs of less than $1M_{\sun}$, one must invoke binarity (mass transfer) to resolve the resulting discrepancy. That supports the scenario of \citet{dm06}, \citet{so06}, and \citet{zi07}, who argue that a large fraction of observable planetaries may indeed stem from binary systems. Consequently, if a planetary nebula/open cluster association is established, we must be aware of the possibility that binarity might negate possible predictions for progenitor mass on the basis of the cluster's implied age from its main sequence turnoff. In this paper we consider the possibility of cluster membership for a number of planetary nebulae that are located in close proximity to open clusters lying in their lines of sight. The often cited cases for planetary nebula/open cluster associations include the cluster and nebula designated as NGC 2818, as well as A9 in NGC 1912 (M38) and NGC 2438 in NGC 2437 (M46), but lesser known cases are also considered. | We have yet to establish a single physical association between a planetary nebula and an open cluster based on a correlation between their radial velocities, reddenings, and distances. However, further follow-up is indicated for a number of cases where the evidence is suggestive, namely M 1-80/Berkeley 57, NGC 2438/NGC 2437, NGC 2452/NGC 2453, VBRC 2 \& NGC 2899/IC 2488, and HeFa 1/NGC 6067, six of the thirteen coincidences considered. Additional good cases may arise from closer examination of some of the other coincidences noted in Table 5, but most of the associated clusters are as yet unstudied, limiting further progress. Almost all potential cluster planetary nebulae lie in cluster coronal regions, typically surrounding open clusters for which limited or no photometric data exist. The fact that very few Galactic open clusters have been studied to the extent that both their nuclear and coronal regions are examined \citep{tu96a} only compounds the situation. Further progress requires not only new studies of our Galaxy's many unstudied clusters, but studies of their coronal regions as well. Spectroscopic observations of potentially-associated planetary nebulae would also be of value. | 7 | 10 | 0710.2900 |
0710 | 0710.2288_arXiv.txt | { VERITAS employs a multi-stage data acquisition chain that extends from the VME readout of custom 500 MS/s flash ADC electronics to the construction of telescope events and ultimately the compilation of information from each telescope into array level data. These systems provide access to the programming of the channel level triggers and the FADCs. They also ensure the proper synchronization of event information across the array and provide the first level of data quality monitoring. Additionally, the data acquisition includes features to handle the readout of special trigger types and to monitor channel scaler rates. In this paper we describe the software and hardware components of the systems and the protocols used to communicate between the VME, telescope, and array levels. We also discuss the performance of the data acquisition for array operations. } \begin{document} | VERITAS \cite{icrc07:Veritas} is an array of 4 imaging Cherenkov telescopes designed to record images of gamma rays impacting the atmosphere. Photo-multiplier signals accompanying an image trigger \cite{icrc07:VeritasL3} are processed using a 3-tiered data acquisition system that operates at the VME crate, telescope, and array levels. The VME data acquisition guides the control and read out of the electronics channels. At the telescope level, the Event Builder combines data from each VME crate into events. At the array level, the Harvester collects events from each telescope and the array trigger and forms the final data product. Figure \ref{fig:daqdiag} shows a schematic of the data flow and communication between the processes. \begin{figure*} \begin{center} \includegraphics [width=0.8\textwidth]{icrc1166_fig1.eps} \end{center} \caption{Schematic of the data transfer and communication/control relationships for the data acquisition systems. Thick black arrows indicate data transfer paths and the protocols used. Thick unfilled arrows indicate communication lines. The analog signals passed between the array trigger and the VME acquisition are included as thin arrows. The dashed line encloses processes repeated for each VME crate and the dash-dotted line encloses those repeated for each telescope. } \label{fig:daqdiag} \end{figure*} \subsection{The VME Data Acquisition} The VME Data Acquisition (VDAQ) serves as the interface to five VME crates that participate in the digitization of the PMT signals and the channel-level triggering for each telescope. Four of these contain 500 MS/s flash ADC modules with a clock-trigger module \cite{icrc03:VeritasFADC} and a fifth serves as an auxiliary crate housing a specialized clock-trigger module and a GPS clock (TTM637VME). The FADC electronics and the constant fraction discriminators (CFDs)\cite{icrc03:VeritasCFD} that produce individual channel triggers are housed on 10-channel 9U VME modules. The required complement of FADC channels to accomodate the 499-pixel camera are distributed among 4 crates of 12 or 13 FADC modules. Each crate is controlled by a VMIVME 7807 Intel Pentium M 1.8 GHz single board computer running Linux. An additional Dolphin PCI mezzanine card provides a communication link using the ANSI/IEEE 1596-1992 Scaleable Coherent Interface (SCI) standard. Each of the VME crates and the telescope Event Builder are connected as nodes of a low-latency, high-throughput network. The configuration used in VERITAS achieves transfer rates of 50 MBytes/s. Data transfers and general communication between VDAQ and the telescope Event Builder are conducted via this interface. Each crate can be accessed via GBit ethernet for starting and halting the acquisition process. A 24-sample (48 ns) FADC trace translates into an event fragment size of 3880 bytes for a crate of 13 modules (130 channels). Event fragments are collected and buffered until the accumulated size reaches 8 MBytes ($\sim$2200 events for FADC crates). The buffer is then shipped to the telescope Event Builder. The transfers from the VME crates happen asynchronously for crates containing different numbers of modules. The typical data rate for a crate is $\sim$780 kBytes/s, which is well below the maximum rates permitted by the SCI and the Event Builder. The chief limitation on the system throughput arises from the data transfer rate over the VME backplane. VDAQ is an event-driven process while the telescope and array acquisition processes are buffered. The telescope deadtime is incurred only at the VME level and is dominated by the size of the crate event fragments. For an array trigger rate of 200 Hz, the deadtime at an individual telescope is about 8.5\%. {\bf Hardware Configuration.} Upon initiation, the VDAQ process for each crate generates a physical map of the modules present by type and, if applicable, a unique board identifier. Each crate provides this map to the Event Builder and uses it internally to procure hardware configuration parameters for each FADC channel and CFD. A variety of programmable features of the FADCs and CFDs are configured \cite{icrc03:VeritasFADC,icrc03:VeritasCFD}. While some parameters are unique to an individual physical component and must be associated by a uniqe identifier, others are properties of the telescope and pixel to which a channel is connected. The configuration settings and component mappings are stored in a MySQL database. This solution suits the distributed nature of the crate acquisition and accomodates the mappings required to properly configure the hardware. An additional benefit is an accurate log of the mappings and settings used. {\bf Trigger Synchronization.} After initialization and configuration, the acquisition processes await signals from the array trigger or commands from the Event Builder. The clock-trigger module in each crate generates a busy level during FADC reads when triggers cannot be accepted. To keep the five-crate system aligned, the busy levels are combined into a single telescope busy level and applied locally as a veto to incoming triggers. This signal is sent to the array trigger \cite{icrc07:VeritasL3} to indicate when the telescope cannot accept triggers. {\bf Trigger Type Handling.} Upon receipt of an array trigger, each crate is passed a serialized event mask that contains an event number and a trigger type code. The mask is encoded into the event fragment via the clock-trigger boards to allow sychronization by the Event Builder. The trigger code is read by VDAQ and can be used to indicate special requests for FADC functions from the array trigger. Out-of-time reads of the FADC memory buffer are used to assess pedestal values periodically during observations. Upon receipt of a pedestal trigger code from the array trigger, VDAQ invokes a dedicated hardware command in the FADCs. Pedestal events are included in the data as normal events distinguished by their type code. {\bf Trigger Rate Measurements.} VDAQ accesses CFD trigger rates via the FADC modules. Scaler counts for each channel are read every 400 events, about once every 2 seconds; this is not often enough to impact the deadtime significantly. The CFD scalers are packed as a specially tagged event and shipped to the Event Builder as part of the regular SCI transfer. Scaler reads are included during normal observations to provide direct diagnostics of channel-level triggering. \subsection{The Telescope Event Builder} Each VERITAS telescope has a dedicated Telescope Event Builder which is responsible for combining the event fragments from each of the five VME crates to produce telescope events; these are then written to local disk and sent via GBit ethernet to the Harvester system. The Telescope Event-Building system is a Dual Intel Xeon Server machine running linux and using a local RAID array. Communications with and control of the Event Builder program is achieved through use of CORBA (Common Object Request Broker Architecture, specifically OmniORB). The Event Building Software is written in C++ and is fully multithreaded, typically containing five threads whilst running: Communications, SCI Buffer Acquisition and Parsing, Event Building, Disk Writer, Network Writer. At the start of each night, the Event Builder queries each VDAQ crate for a map of present VME modules and then dynamically configures itself. Data is buffered at each of the VDAQ machines and transferred in blocks to the Event Builder via the SCI system. The Event Builder parses these memory blocks and extracts pieces of individual events, tagged by unique event numbers, which are then stored in memory. When all of the pieces of an event have arrived, the telescope event is built and buffered. Once roughly 160 kBytes of telescope events have been accumulated, the events are then dispatched to the ``Consumers''. The Consumers are processes that receive built data buffers. They utilize a common architecture and, at run-time, any number of consumers may be registered. Typically only two are; the Disk Writer and Network Writer, but an additional Data Integrity Monitor consumer may also be used. It is estimated that the throughput of the Telescope Event-Building system on a dual 2.4 GHz Xeon server is approximately 12 MBytes/s. In addition to receiving actual event data, the Telescope Event Builders periodically receive CFD scaler data as described in the previous section. These data are not transferred with the event data, but simply stored and made available via CORBA to any system that requests the information. \subsection{The Harvester} VERITAS back-end data acquisition is the task of the Harvester -- a single eight-core machine that collects data from all telescopes in real-time in addition to a stream of meta-data from the L3 trigger. The Harvester accomplishes four tasks: {\bf Storing data.} The Harvester saves all data streams to its fast RAID in real time. The current strategy for real-time data storage is to create a separate file for each telescope; this makes it possible to handle the separate telescope streams in parallel with minimal interaction. {\bf Combining data.} In addition to saving the data, the Harvester combines it in real time into \emph{array events}. In this way, real time diagnostics can see a ``big picture'' of what the array is doing, rather than having to deal with telescopes individually. {\bf Diagnostics.} The Harvester runs a variety of real-time diagnostics -- ranging from sanity checks to see if telescopes read out when they were supposed to, to a complete high-performance stereo analysis package that serves as the VERITAS real-time quicklook. During a run, the observer has up-to-the-second knowledge of the performance of the system. {\bf Creating the final product.} After a run completes, the Harvester immediately starts combining the data streams from the run to create a single file using the VERITAS Bank File (VBF) data format. \smallskip VBF groups telescope events together, such that given an event number, the user has immediate access to the corresponding events from all telescopes, in addition to meta-data from the array trigger. Thus, the analysis does not need to correlate stereo events; this task is already accomplished by the data acquisition system. VBF has been designed for portability, high performance reading and writing, compactness, and extensibility. High performance access and compactness are achieved using a custom scheme for compressing FADC samples based on picking a different number of bits-per-sample depending on the dynamic range of each particular trace. This scheme overwhelmingly outperforms gzip and bzip2 in reading and writing times while reaching similar compactness. See Table~\ref{vbfperf} for performance measurements \begin{table} \begin{center} \begin{tabular}{|l|c|c|} \hline VBF Variant & Space Usage & Read Time \\ \hline \hline Uncompressed & 100\% & 100\% \\ % w/ Gzip & 42\% & 114\% \\ % w/ Bzip2 & 35\% & 514\% \\ \hline Compressed$^*$ & 38\% & 64\% \\ % Comp. w/ Gzip & 32\% & 93\% \\ % Comp. w/ Bzip2 & 30\% & 471\% \\ \hline \end{tabular} \end{center} \caption{Comparison of compression schemes normalized to the uncompressed VBF case. performed on a Pentium 4 with 2 GB RAM, an eight-way SCSI RAID-0, and Linux 2.4.18. For bzip2, we used a block size of 100,000 bytes. $^*$VERITAS uses custom compression alone. } \label{vbfperf} \end{table} The observer interacts with the Harvester using the VERITAS array control software, as well as a suite of GUIs designed to view the results of quicklook analysis. The stereo analysis performed by the VERITAS quicklook system is capable of a very high level of performance -- both in time and in sensitivity. A typical twenty minute run takes two minutes to analyze using quicklook. Further, the quicklook analysis results are comparable to offline analysis packages. Thus, we have confidence that if a bright source appeared in our field of view during observations, quicklook would have as good of a chance of seeing it as a manually performed offline analysis. | The VERITAS data acquisition systems combine a variety of hardware and software resources to achieve efficient and reliable operation from the reading of FADC data to the building and storage of event data. Diagnostic information is available at all levels and real-time analsyis is performed to ensure data quality. The system has proven highly adaptable and meets the needs of a variety of configuration and calibration tasks. All systems operate within design parameters and leave room for the exploration of low-threshold regimes. \subsection* | 7 | 10 | 0710.2288 |
0710 | 0710.2241_arXiv.txt | {} {We present a uniform catalog of the images and radial profiles of the temperature, abundance, and brightness for 70 clusters of galaxies observed by {\it XMM-Newton}.} {We use a new ``first principles'' approach to the modeling and removal of the background components; the quiescent particle background, the cosmic diffuse emission, the soft proton contamination, and the solar wind charge exchange emission. Each of the background components demonstrate significant spectral variability, several have spatial distributions that are not described by the photon vignetting function, and all except for the cosmic diffuse emission are temporally variable. Because these backgrounds strongly affect the analysis of low surface brightness objects, we provide a detailed description our methods of identification, characterization, and removal.} {We have applied these methods to a large collection of XMM-Newton observations of clusters of galaxies and present the resulting catalog. We find significant systematic differences between the {\it Chandra} {\rm and} {\it XMM-Newton} {\rm temperatures.}} {} | Clusters of galaxies are the largest and most massive collapsed objects in the universe, and as such they are sensitive probes of the history of structure formation. While first discovered in the optical band in the 1930s \citep[for a review see][]{b97}, in some ways the name is a misrepresentation since most of the baryons and metals are in the diffuse hot X-ray emitting intercluster medium and not in the galaxies. Clusters are fundamentally ``X-ray objects'' as it is this energy range where this preponderance of the baryons is visible. Studies of cluster evolution can place strong constraints on all theories of large scale structure and determine precise values for many of the cosmological parameters. As opposed to galaxies, clusters probably retain all the enriched material created in them, and being essentially closed boxes they provide a record of nucleosynthesis in the universe. Thus measurement of the elemental abundances and their evolution with redshift provides fundamental data for the origin of the elements. The distribution of the elements in clusters reveals how the metals moved from stellar systems into the IGM. Clusters should be fair samples of the universe and studies of their mass and their baryon fraction should reveal the gross properties of the universe as a whole. Since most of the baryons are in the gaseous phase and clusters are dark-matter dominated, the detailed physics of cooling and star formation are much less important than in galaxies. For this reason, clusters are much more amenable to detailed simulation than galaxies or other systems in which star formation has been a dominant process. Clusters are luminous, extended X-ray sources and are visible out to high redshifts with current technology. The virial temperature of most groups and clusters corresponds to $T\sim2-100\times10^6$~K ($kT\sim0.2-10$~keV, velocity dispersions of $180-1200$~km~s$^{-1}$), and while lower mass systems certainly exist we usually call them galaxies. Most of the baryons in groups and clusters of galaxies lie in the hot X-ray emitting gas that is in rough virial equilibrium with the dark matter potential well \citep[the ratio of gas to stellar mass is $\sim2-10:1$,][]{asf01}. This gas is enriched in heavy elements \citep{mea78} and it thus preserves a record of the entire history of stellar evolution in these systems. The presence of heavy elements is revealed by line emission from H and He-like transitions as well as L-shell transitions of the abundant elements. Most clusters are too hot to have significant ($>2$~eV equivalent width) line emission from C or N, although cooler groups may have detectable emission from these elements. However, all abundant elements with $z>8$ (oxygen) have strong lines from H and He-like states in the X-ray band and their abundances can be well determined. Clusters of galaxies were discovered as X-ray sources in the late 1960's (see \citep[for a historical review see][]{m02} and large samples were first obtained with the {\it Uhuru} satellite in the early 1970's \citep{jf78}. Large samples of X-ray spectra and images were first obtained in the late 1970's with the {\it HEAO} satellites \citep[for an early review see][]{jf84}. The early 1990's brought large samples of high quality images with the {\it ROSAT} satellite and good quality spectra with {\it ASCA} and {\it Beppo-SAX}. In the last few years there has been an enormous increase in the capabilities of X-ray instrumentation due to the launch and operation of {\it Chandra} and {\it XMM-Newton}. Both {\it Chandra} and {\it XMM-Newton} can find and identify clusters out to $z>1.2$ and their morphologies can be clearly discerned to $z>0.8$. Their temperatures can be measured to $z\sim1.2$ and {\it XMM-Newton} can determine their overall chemical abundances to $z\sim1$ with a sufficiently long exposure. For example, a cluster at $z=1.15$ has recently had its temperature and abundance well constrained by a 1~Ms {\it XMM-Newton} exposure \citep{hea04}. The temperature and abundance profiles of clusters out to redshifts of $z\sim0.8$ can be measured and large samples of X-ray selected clusters can be derived. {\it Chandra} can observe correlated radio/X-ray structure out to $z>0.1$ and has discovered internal structure in clusters. The {\it XMM-Newton} grating spectra can determine accurate abundances for the central regions of clusters in a model independent fashion for oxygen, neon, magnesium, iron, and silicon. Despite the stunning successes of the {\it Chandra/XMM-Newton} era, clusters have not yet fulfilled their promise as a cosmological Rosetta stone; the most important tests of cluster theory require measurements of cluster properties to large radii ($R\sim R_{virial}$) which is observationally difficult. The lack of consensus among the recent X-ray missions about, for example, temperature profiles, is a large stumbling block in the use of clusters for cosmological purposes. \subsection{Temperature Structure of Clusters} As discussed in detail by \citet{e03}, we now have a detailed understanding of the formation of the dark matter structure for clusters of galaxies. If gravity has been the only important physical effect since the formation, then the gas should be in rough hydrostatic equilibrium and its density and temperature structure should provide a detailed measurement of the dark matter distribution in the cluster. Recent theoretical work has also taken into account other processes, such as cooling, which can be important. The fundamental form of the \citet{nfw97} dark matter potential and the assumption that the fraction of the total mass that is in gas is constant with radius results in a prediction that the cluster gas should have a declining temperature profile at a sufficiently large distance from the center (in $R/R_{viral}$ units), both from analytic \citep{ks01} and numerical modeling \citep{lea02}. The size of the temperature drop in the outer regions is predicted to be roughly a factor of 2 by $R/R_{viral}\sim0.5$. Although some observational results appear consistent with the theoretical predictions \citep[in particular,the {\it ASCA} results of ][]{mea98}, many others do not, and considerable controversy exists. Much of the uncertainty of the pre-{\it Chandra}/pre-{\it XMM-Newton} data arises from insufficient spectral and spatial resolution and the resultant difficulties in removing backgrounds, modeling the spectra, and interpreting the results. For example, the {\it ASCA} results of \citet{mea98} were consistent with a decline in temperature with radius, while the analysis of a similar sample of clusters by \citet{kea99}, \citet{wb00}, and \citet{w00} revealed a large number of isothermal clusters. Similar results were obtained from {\it Beppo-SAX}, with \citet{dgea99} finding temperature gradients and \citet{ib00} finding isothermality. Simultaneous analysis of the higher angular resolution {\it ROSAT} data with the {\it ASCA} data did not resolve the issue; \citet{fad01} finding gradients and \citet{ibe99} isothermal profiles. The bulk of the problem with interpreting {\it ASCA} results is the analysis of impact of the PSF on the profile \citep{ibe99}. {\it XMM-Newton} and {\it Chandra} have significantly better spectral and angular resolution than the previous generation of missions and might be expected to resolve the previous controversies. The recent {\it Chandra} results of \citet{vea06} show a temperature profile in good agreement with the gradients seen by \citet{mea98} results and predicted by the standard theory. Analysis of samples of cooling flow clusters with {\it XMM-Newton} \citep{pea05,app05,pea07} are also mostly consistent with the \citet{mea98} results. However flatter, more isothermal profiles have also been found in both {\it Chandra} and {\it XMM-Newton} observations \citep{asf01,kea04,app05}. Despite some early difficulties \citep[e.g.,][]{dea06}, the {\it Chandra} and {\it XMM-Newton} calibrations have stabilized but agreement between the two great observatories is not assured \citep[e.g,][]{vea06}. The difference in the PSF between the two instruments as well as different methods of background subtraction often make direct comparison difficult. Further, an agreement between {\it Chandra} and {\it XMM-Newton} would not entirely resolve the problem; the smaller FOV of current instruments have led to observation of a somewhat higher redshift sample than observed by the previous generation of instruments, suggesting that part of the difference between the {\it XMM-Newton}/{\it Chandra} results and the {\it ASCA}/{\it ROSAT}/{\it Beppo-SAX} results may be due to a real difference between clusters at lower and higher redshifts. The measurement of the cluster mass function can provide a sensitive cosmological test but is sensitive, in turn, to the parameters that are directly measurable, and especially to the observed quantities at large radius. Recent simulations show that cluster temperature profiles decline with radius but less rapidly than is shown by previous X-ray analysis \citep[e.g.,][]{ktea04}. Since the total mass of the cluster is quite sensitive to the measured temperature profile \citep{rea06}, particularly at large radii, these systematic differences lead to significant uncertainties in the cosmological constraints. Thus, there is an urgent need to understand the temperature profiles of clusters at large radii and to understand the source of the systematic differences observed in the literature. In this paper we consider a large sample of clusters observed with the {\it XMM-Newton} observatory and derive temperature, density and abundance profiles for many of these systems out to near the virial radius. We present a new technique that should provide more accurate background subtraction at large radii, and are careful to correct for the effect of the finite {\it XMM-Newton} PSF. A comparison of our measurements with {\it Chandra} measurements of the same clusters shows a simple systematic difference between the two observatories. Although we have not yet determined the source of that difference, resolution of this relatively well defined issue should significantly reduce the uncertainties in cluster cosmology. \subsection{Analysis of Extended Sources} The analysis of extended sources in X-ray astronomy is typically problematic and quite often very complex. This is particularly true for objects which subtend the entire field of view (FOV) of the observing instrument such as nearby galaxies, relatively nearby clusters of galaxies, many regions of galactic emission, and of course the cosmic diffuse background. Even with objects smaller than the FOV, quite often the simple subtraction of a nearby ``background'' region from the same data set is inappropriate due to spectral and spatial variations in the internal background and angular variations in the cosmic background. The use of deep ``blank sky'' observations can also be inappropriate due to the same considerations, as well as the probability that many background components are temporally varying. Because of the temporal variation of the background and the angular variation of the cosmic background, the average of the blank-sky data, even after normalization, may not match the conditions of a specific observation of interest, and so may yield an incorrect result. While the cores of many clusters are relatively bright in X-rays so the treatment of the background is not such a significant consideration, at the edges of clusters it is absolutely critical. Clusters fade gently into the backgrounds at large radii, therefore improving the modeling of the backgrounds extends the reliable radial range for the determination of cluster parameters. Critical to compensating for the various background components by filtering, subtraction, or modeling is a basic understanding of their origin and effects on the detectors. Unfortunately this usually takes a considerable amount of time to develop, which is why useful methods for a specific observatory become available to the general community only years into the mission. Even then, the methods will continue to evolve with greater understanding of the various background components and their temporal evolution, and the operation of the instruments. In addition, the efforts are quite often undertaken by individuals who are not project personnel, but whose scientific interests require the improved analysis methods. This is certainly true of the {\it XMM-Newton} mission and observations using the EPIC instruments. Several groups have presented methods and published scientific results based upon them \citep{aea01,rp03,nml05,dlm04}. As opposed to these methods which derive background spectra from normalized blank-sky observations, this paper presents the details of a method based as much as possible on the specific understanding of the individual background components. This method was used successfully in the paper identifying the solar wind charge exchange emission in the {\it XMM-Newton} observation of the Hubble Deep Field North \citep{sck04}. Section~\ref{sec:instrumentation} of this paper gives a short description of the {\it XMM-Newton} observatory, Sect.~\ref{sec:components} discusses the various background components and the suggested methods used to compensate for them, Sect.~\ref{sec:example} demonstrates the data reduction method using the observation of \object{Abell 1795}. Sect.~\ref{sec:catalog} applies the methods to the determination of the temperature, abundance, and flux radial profiles of a catalog of 70 clusters of galaxies and presents the results, and Sect.~\ref{sec:conclusions} discusses the conclusions. Note that the detailed discussion of the science derived from these observations is deferred to Paper~II. Currently the specific method and software package discussed here are only applicable to EPIC MOS data. Although the MOS and pn experience the same backgrounds, the physical difference between the two detectors (readout rates, fraction of unexposed pixels, etc.) make analysis of the pn background somewhat more difficult than that of the MOS. However, the analysis methods described here are being extended to the pn. | \label{sec:conclusions} In this paper we have outlined a robust and reliable method for analyzing extended X-ray sources observed with the {\it XMM-Newton} EPIC MOS detectors. The method combines screening of the data for periods of background enhancements (most notably the soft proton contamination), detailed modeling of the particle background spectrum, and the determination of other background components in the spectral fitting process (residual SP contamination, fluorescent particle background lines, and the cosmic background). We have demonstrated our method with the bulk processing of the observation of 70 clusters of galaxies. Comparison of the results for two separate observations of \object{Abell 1835}, \object{S\'ersic 159-3}, and \object{Perseus} show good agreement between their fitted temperatures. However, comparison of our results with the {\it Chandra} results of \citet{vea05} for the overlapping subset of clusters shows a significant discrepancy for higher temperature clusters. The sense of this discrepancy is that the higher the fitted temperature, the greater the likelihood that {\it Chandra} will find a higher temperature than {\it XMM-Newton}. The differences can be over 1~keV at $7-8$~keV. This effect can increase the apparent temperature gradient in the outer annuli of clusters in {\it Chandra} data. While the detailed scientific analysis and discussion of these results are deferred to Paper~II, a few aspects are clear from plots of the entire data set. For the combined plots, the radii of the annuli have been scaled to the $R_{500}$ value in the same manner as the individual plots (Sect.~\ref{sec:catalog}). Figs.~\ref{fig:temp}, \ref{fig:abund}, and \ref{fig:flux} show the cumulative plots for the temperature, abundance, and flux, respectively. Again, both the temperature and flux have also been normalized to the values in the range 5\% -- 30\% of $R_{500}$. In addition, only points where the fitted values are three times the fitted uncertainty are plotted. \begin{figure} \centering{\includegraphics[angle=0,width=8.5cm]{7930-18-temp.pdf}} \caption{Scaled temperature radial profiles for all of the analyzed clusters. \label{fig:temp}} \end{figure} Inspection of Fig.~\ref{fig:temp} shows, as seen before \citep[e.g.][]{pea07,app05,vea06}, a wide variety of temperature profiles inside 5\% of R(500). Most of these can be characterized by a temperature drop in the center as has long been seen in cooling-flow clusters. However our single phase analysis may produce results slightly different than more detailed analysis. Over the range from $0.05-0.2 R_{500}$ the clusters are isothermal to better than 5\%. Beyond $\sim0.2 R_{500}$ a significant fraction of the clusters (Paper~II) show temperature drops, but they are not all self-similar. However a significant fraction of the clusters are relatively isothermal out to the largest radii measurable. As noted by \citet{app05}, many of the clusters show a self-similar surface brightness profile (Fig.~\ref{fig:flux}). Inside of $\sim0.03 R_{500}$ there is significant scatter in the profile. With respect to the overall abundance, as was noted for {\it ASCA} spectra of clusters by \citet{fad01} and later for many {\it XMM-Newton} and {\it Chandra} spectra \citep{mea07} there is, in a significant fraction of the clusters, an abundance increase in the center. However outside of the central $\sim0.05 R_{500}$ there is little evidence for an abundance gradient and all the clusters are very close to the average value of $A=0.3$ on the \citet{ag89} abundance scale (Fig.~\ref{fig:abund}). Detailed analysis of these results will appear in Paper~II. \begin{figure} \centering{\includegraphics[angle=0,width=8.5cm]{7930-19-abund.pdf}} \caption{Abundance radial profiles for all of the analyzed clusters. \label{fig:abund}} \end{figure} \begin{figure} \centering{\includegraphics[angle=0,width=8.5cm]{7930-20-flux.pdf}} \caption{Scaled flux radial profiles for all of the analyzed clusters. \label{fig:flux}} \end{figure} | 7 | 10 | 0710.2241 |
0710 | 0710.0558_arXiv.txt | Deep F555W and F814W {\it Hubble Space Telescope\/} ACS images are the basis for a study of the present day mass function (PDMF) of NGC~346, the largest active star forming region in the Small Magellanic Cloud (SMC). We find a PDMF slope of $\Gamma=-1.43\pm 0.18$ in the mass range $0.8-60\, {\rm M}_\odot$, in excellent agreement with the Salpeter Initial Mass Function (IMF) in the solar neighborhood. Caveats on the conversion of the PDMF to the IMF are discussed. The PDMF slope changes, as a function of the radial distance from the center of the NGC~346 star cluster, indicating a segregation of the most massive stars. This segregation is likely primordial considering the young age ($\sim 3\, {\rm Myr}$) of NGC~346, and its clumpy structure which suggests that the cluster has likely not had sufficient time to relax. Comparing our results for NGC~346 with those derived for other star clusters in the SMC and the Milky Way (MW), we conclude that, while the star formation process might depend on the local cloud conditions, the IMF does not seem to be affected by general environmental effects such as galaxy type, metallicity, and dust content. | \label{intro} Since the pioneering study by \citet{salpeter55}, one of the major issues in astrophysics is whether the initial mass function (IMF), the relationship that specifies the mass distribution of a newly formed stellar population, is universal or, alternatively, determined by environmental effects. In particular, establishing whether the IMF has been constant over the evolution of the universe, or if it varies with time (redshift) and/or metallicity has crucial consequences on the evolution, surface brightness, chemical enrichment and baryonic content of galaxies, and on the evolution of light and matter in the universe. Due to a number of effects and observational uncertainties such as mass segregation, field star contamination, and variable extinction \citep{elmegreen06}, the ``true" IMF is very difficult to measure. Despite these limitations, star clusters still provide the best determination of the IMF, because they are essentially single--age, single--metallicity stellar systems. In general, all star clusters lose more than 80\% of their stellar content in their first $\sim 10$ Myr \citep{kroupa02,lada03}. Therefore, in order to perform an accurate census of the initial stellar content it is necessary to focus on the youngest stellar systems. In addition, young (age $<3-5\, {\rm Myr}$) star clusters in nearby galaxies (e.g. the Large --LMC-- and Small Magellanic Cloud --SMC) offer to us the unique advantage to individually resolve --- because of their proximity --- statistically significant samples of stars down to the sub--solar mass regime. In this paper, we focus our attention on the mass function (MF) of NGC 346, an extremely young \citep[$\sim 3\, {\rm Myr}$,][hereafter S07]{bouret03,sabbi07}, and active star cluster that excites the largest and brightest H{\sc ii} region (N66) in the SMC. NGC~346 ($\alpha_{j2000}=00^h 59^m 05.2^s, \delta_{j2000}=-72^{\deg} 10'28''$) contains a major fraction of the O stars known in the entire SMC \citep{walborn78,walborn86,niemela86,massey89}. The bright end of its stellar population ($V\le 19.5\, {\rm mag}$) has been well investigated in the past 20 years. Spectral investigations of the brightest members identified several stars of spectral type O6.5 or earlier \citep{massey89,heap06,mokiem06,evans06}. \citet{massey89} found the MF for the brightest stars in NGC 346 --- down to 5 M$_{\odot}$ --- to have a slope that is consistent with what has been found for massive stars near the Sun and in the LMC. High resolution observations obtained with the Advanced Camera for Survey (ACS) on board of the \emph{Hubble Space Telescope} (\emph{HST}) recently revealed that the star formation history of the region where NGC~346 is located is quite complex. There is evidence for old episodes of star formation (SF), in the field of the SMC approximately between 3 and 5 Gyr ago. After that, the SF activity in this region appears to have been significantly diminished, with a possible moderate enhancement $\sim 150\, {\rm Myr}$ ago (S07). An intermediate--age ($\sim 4.5\, {\rm Gyr}$) stellar cluster (BS90), characterized by a core radius $r_c\simeq 25\, {\rm arcsec}$ and a tidal radius $r_t \simeq 130\, {\rm arcsec}$, is located at a projected distance of $\sim 23\, {\rm pc}$ northeast from the center of NGC~346 (S07). Spectral analysis of NGC~346 members revealed the presence of stars as young as $\sim 1\, {\rm Myr}$, \citep{massey05}, but stars $\sim 5\, {\rm Myr}$ old are also present \citep{heap06,mokiem06}. Star formation is likely still active in the region: hundreds of pre--main sequence (pre--MS) stars in the mass range $0.6 - 3\, {\rm M}_{\odot}$ were discovered by \citet{nota06} from \emph{HST}/ACS images. \emph{Spitzer Space Telescope} (\emph{SST}) observations also detected a myriad of embedded young stellar objects (YSOs) scattered across the entire region \citep{bolatto07}. By assuming a Salpeter IMF, \citet{simon07} calculated that more than 3000 M$_\odot$ have been formed in the last $\sim 10^6\, {\rm yr}$, concluding that $\ge 6$\% of the current SF in the SMC is taking place in NGC~346. This recent stellar population does not appear to be uniformly distributed within the ionized nebula, but is rather organized in many, likely coeval, subclusters (S07) that are coincident with clumps of molecular gas previously detected by \citet{rubio00}. \citet{simon07} noted that almost all the subclusters host one or more YSOs. The high sensitivity of the \emph{HST}/ACS allowed us to investigate the stellar content of NGC~346 from $\sim 60 M_\odot$ down to $\sim 0.6 M_\odot$, making NGC~346 one of the few known regions where a MF can be determined over two orders of magnitude \citep[the other classical case being the Orion trapezium cluster in our own MW;][]{elmegreen06}. This paper gives a new derivation of the NGC 346 MF, reviews the strengths and limitations of the methods, and presents the conclusions that we can conservatively derive. The paper is organized as follows: a short description of the data reduction procedure is presented in \S\ref{obs}; in \S\ref{CMD} we present the color--magnitude diagram (CMD), and we describe the stellar populations identified. In \S\ref{MF} we present the MF we obtained for the NGC~346 cluster, we discuss its spatial variations, and we analyze the impact of the environment on the MF. The results of this paper are discussed in \S\ref{conclusions}. | \label{conclusions} Our analysis of the stellar content of NGC~346 indicates that, within all the uncertainties that can affect the determination of a MF, the PDMF slope of NGC 346, at least down to the solar mass regime, is consistent with the value derived by \citet{salpeter55} for the IMF in the solar neighborhood, further confirming that the stellar MF does not strongly depend on the metallicity, or the mass or morphological type of the parent galaxy. We studied how the PDMF varies as a function of the distance from the center of NGC~346. We found that the PDMF is quite flat in the center, and becomes steeper in the periphery. As already found in other young star clusters \citep{hillenbrand97,hillenbrand98,sirianni02} the increase in the MF slope with the distance from the center is due to the lack of massive stars in the periphery rather than an excess of low--mass objects there. The projected density of massive stars from the innermost annulus and the outskirts decreases by a factor of 60, whereas low-mass stars are depleted only by a factor of 6; the most massive stars are segregated in the center. We conclude that the spatial stellar mass distribution is unlikely to be biased by age effects. SST observations indicate that there is ongoing SF throughout the complex with a lower limit on the total star formation rate of $3.2 \times 10^{-3}\, {\rm M}_\odot\, {\rm yr}^{-1}$ \citep{simon07}. Furthermore \citet{mokiem06} noted a spread of about 5 Myr in the ages of NGC~346 stars, but did not find any correlation between the age and the position of the stars over the whole cluster. S07 noted that stars in NGC~346 appear to be organized in several coeval sub--clusters, embedded in H{\sc ii} gas, and coinciding with the CO clumps analyzed by \citet{rubio00}. YSOs between Class I and Class III are found in 14 of the 16 identified sub--cluster \citep{simon07}. The sub--clusters are also connected by filaments and arcs of gas and dust. The complex structure of NGC~346 and its young age, compared to its dynamical time, suggest that NGC~346 is not dynamically evolved, and probably not yet relaxed. On the basis of these considerations we conclude that the observed segregation of the most massive stars is likely primordial, and reflects how the cluster formed. S07 suggested that in NGC~346 most of the star formation was likely not triggered, but rather resulted from the turbulence driven density variations within a giant interstellar cloud complex, conditions predicted by the hierarchical fragmentation of a turbulent molecular cloud model \citep{elmegreen00,klessen00,bonnell02,bonnell03}. According to this model, the fragmentation of the cloud is due to supersonic turbulent motions present in the gas. The turbulence induces the formation of shocks in the gas, and produces filamentary structures \citep{bate03}. The chaotic nature of the turbulence increases locally the density in the filamentary structures. When regions of high density become self--gravitating, they start to collapse to form stars. Simulations show that star formation occurs simultaneously at several different locations in the cloud \citep{bonnell03}, as appears to be the case for NGC~346. In this scenario, stars forming near the cluster center would have higher accretion rates, due to the local higher gas density \citep{larson82,bonnell01}. Simulations by \citet{bonnell01} indicate that although low--mass stars form equally well throughout the entire cluster, the most massive stars are almost exclusively segregated to the central regions. It is interesting to note that the two most massive stars (one of them being an O3V spectral type, Figs.~\ref{f:centro}) are $\sim 7\, {\rm pc}$ off the center of the cluster (outside the core of NGC~346, but still within the half mass radius). Did these stars form where they are observed today, in a quite low stellar density region? \citet{vine03} studied the effect of strong stellar winds and photoionizing fluxes coming from O- and B-type stars on the gas content in a young star cluster. According to their simulations, the gas expulsion is unlikely to remove any initial mass segregation, but it can still result in the ejection of some of the most massive stars to positions outside the cluster core, although still within the half mass radius. According to \citet{bouret03} mechanical power from O-star winds is not a dominant factor in the evolution within N66, but the molecular gas in the original cloud has been strongly photodissociated and photoionized by the O stars \citep{rubio00}. \citet{niemela86} measured the radial velocity of five of the innermost and brightest stars in NGC~346, and found a velocity dispersion of $\sim 5\, {\rm km\, s}^{-1}$, indicating that, even in $\sim 1\, {\rm Myr}$, they can be formed in the center of the cluster and then ejected where they are observed today. Another possibility is that these stars formed where are observed, as the result of a very efficient competitive accretion process into a potential well \citep{bonnell03}. Further spectroscopic analysis of the dynamics of the gas and the sub--clusters can confirm whether the giant molecular cloud which formed NGC~346 underwent under a hierarchical fragmentation. The initial supersonic turbulence originates shocks in the gas which rapidly remove kinetic energy from the gas \citep{ostriker01}: in this case we expect to find low gas velocities around the sub--clusters. A detailed study of the sub--clusters dynamics will also reveal if NGC~346 is or is not already relaxed and virialized. We will present these results in a future paper. | 7 | 10 | 0710.0558 |
0710 | 0710.0591_arXiv.txt | We report the first fully sampled maps of the distribution of interstellar CO$_2$ ices, H$_2$O ices and total hydrogen nuclei, as inferred from the 9.7 $\mu$m silicate feature, toward the star-forming region Cepheus A East with the IRS instrument onboard the {\it Spitzer Space Telescope}. We find that the column density distributions for these solid state features all peak at, and are distributed around, the location of HW2, the protostar believed to power one of the outflows observed in this star-forming region. A correlation between the column density distributions of CO$_2$ and water ice with that of total hydrogen indicates that the solid state features we mapped mostly arise from the same molecular clumps along the probed sight lines. We therefore derive average CO$_2$ ice and water ice abundances with respect to the total hydrogen column density of $X$(CO$_2$)$_{ice}$$\sim$1.9$\times$10$^{-5}$ and $X$(H$_2$O)$_{ice}$$\sim$7.5$\times$10$^{-5}$. Within errors, the abundances for both ices are relatively constant over the mapped region exhibiting both ice absorptions. The fraction of CO$_2$ ice with respect to H$_2$O ice is also relatively constant at a value of 22\% over that mapped region. A clear triple-peaked structure is seen in the CO$_2$ ice profiles. Fits to those profiles using current laboratory ice analogs suggest the presence of both a low-temperature polar ice mixture and a high-temperature methanol-rich ice mixture along the probed sightlines. Our results further indicate that thermal processing of these ices occurred throughout the sampled region. | Understanding how the structure and composition of ice mantles covering interstellar grains change with the local physical and chemical conditions is crucial to constraining the chemical evolution of protostellar envelopes, protoplanetary disks, and comets. Infrared observations toward low-to-high mass star-forming regions (e.g., Nummelin et al. 2001; Gibb et al. 2004) and quiescent molecular clouds (Whittet et al. 1998) using the {\it Infrared Space Observatory} have shown that water (H$_2$O), carbon monoxide (CO) and carbon dioxide (CO$_2$) are common constituents of the ice mantles covering interstellar grains in dense molecular environments. While H$_2$O and CO are believed to form from grain surface reactions and gas-phase desorption, respectively (e.g., d'Hendecourt et al. 1985), the production of solid CO$_2$ remains uncertain. Mechanisms such as UV photolysis and/or grain surface chemistry are invoked. The routine detection of solid CO$_2$, with a fraction relative water ice between 9 and 37 \% toward a number of star-forming regions (e.g., Gerakines et al. 1999; Nummelin et al. 2001), combined with the fact that CO$_2$ is easily produced via UV irradiation in the laboratory (Ehrenfreund et al. 1996), added weight to the suggestion that CO$_2$ ices are mainly produced via UV photolysis in the interstellar medium. However, the detection of CO$_2$ ice toward a number of quiescent clouds devoid of any embedded source of radiation, with abundances relative to water ice similar to those measured toward star-forming regions, demonstrated that production mechanisms such as grain surface reactions and gas-grain interaction need be considered too (Whittet et al. 1998; Whittet et al 2007). A number of laboratory experiments studied the evolution of ice mixtures when subjected to thermal and/or UV processing (e.g., Ehrenfreund et al. 1996; 1999). Comparisons of these laboratory ice analogs with {\it ISO} and more recent {\it Spitzer} observations of interstellar ices allowed previous workers to better constrain the physical and chemical conditions in the clouds toward which these ices were detected (e.g., Ehrenfreund et al. 1998; 1999). In particular, the profile of the $\nu_2$ vibrational bending mode of solid CO$_2$ (15.2 $\mu$m) was found to be very sensitive to the composition of the ice mantle the molecules are embedded in and to the thermal history of the region under study. Substructures that are characteristic of ice crystallization and segregation of the ice mantle constituents upon thermal processing have been routinely used to constrain the physical and chemical environments around young stellar objects (e.g., Boogert et al. 2000; Gerakines et al. 1999; Boonman et al. 2003; Gibb et al. 2004; Bergin et al. 2005). These previous studies were all based on pointed observations of individual discrete sources, and thus require observations of multiple sources to provide any spatial information. With the {\it Spitzer} spectral mapping capability, one can now investigate the distribution of interstellar ices within a given dense cloud or star-forming region with unprecedented spatial sampling ($\sim$3$''$). This provides a new opportunity to link the spectral properties of the ices with local physical conditions, and thus to throw new light on their origin and evolution. In this paper, we report on the spatial distribution and the evolution of the interstellar ices present toward the star-forming region Cepheus A East using this new capability offered by {\it Spitzer}. Cepheus A is a well-known site of star formation located at a distance of about 650 pc. The region contains a series of deeply embedded far-infrared and radio-continuum sources, one of which dominates the luminosity of the entire region (HW2 with 2.5$\times$10$^{4}$ $L_{\odot}$; Hughes \& Wouterloot 1984; Lenzen et al. 1984; Evans et al. 1981). Ground- (e.g., Hartigan et al. 1996; Goetz et al. 1998) and space-based observations (e.g., Wright et al. 1996; van den Ancker et al. 2000) revealed the existence of a multipolar outflow exhibiting complex structures of shock-excited atomic and molecular gas components, indicating the presence of both dissociative (J-type) and non-dissociative (C-type) shocks resulting from successive episodes of activity (Narayanan \& Walker 1996). The protostellar object HW2 was found to be the dominant powering source of the extremely high-velocity (EHV) outflow oriented {\it northeast} ({\it NE}), while the source (or sources) of the high-velocity jet (HV; oriented {\it southeast}) is still under debate (e.g., Goetz et al. 1998; Hiriart et al. 2004; Codella et al. 2003; Mart\'{\i}n-Pintado et al. 2005). Previous observations using {\it ISO} exhibited a number of absorption bands due to the presence of ice mantles containing CO, CO$_2$ and H$_2$O and due to silicate grains over the {\it NE} side of this star forming region (van den Ancker et al. 2000). The CO and CO$_2$ ice abundances -- with respect to water ice and averaged over the Short-Wavelength Spectrometer (SWS) aperture -- were found to be within the range observed toward other sight lines (e.g., Whittet et al. 1996). Using new high-resolution {\it Spitzer} data, we recently reported the first detection of gas-phase CO$_2$ emission. This emission extends over a 35$''\times$25$''$ region associated with the EHV outflow (Sonnentrucker et al. 2006). We determined that the gaseous CO$_2$ molecules mostly result from sputtering of ice mantles due to interactions between the EHV outflow and the ambient molecular medium. We derived a CO$_2$ gas-to-ice ratio of at most 3\% over the region showing CO$_2$ gas emission. Our data also exhibit strong absorption from the 15 $\mu$m CO$_2$ ice bending mode, from the 6 $\mu$m water ice band and from the 9.7 $\mu$m silicate feature at the position of HW2 and at numerous spatial positions preferentially located away from the EHV outflow region. In this paper, we discuss the distribution of the solid state features observed toward Cepheus A East with the {\it Spitzer} Infrared Spectrograph (IRS). While H$_2$ $S$(0) to $S$(7), [NeII], [NeIII], [S I], [SIII] and [FeII] emissions as well as emission from C$_2$H$_2$ (Sonnentrucker et al., submitted) are also detected in our data, we focus here on the study of the absorptions from CO$_2$ ices, H$_2$O ices and silicate grains -- the latter used to estimate the total column density of hydrogen nuclei $N$(H$_{\rm{tot}}$) -- toward Cepheus A East. Section 2 summarizes our observations and Section 3 describes our data analysis. Section 4 compares the spatial distribution of the CO$_2$ ice, H$_2$O ice and silicate grains with that of the quiescent molecular clouds present in this region. Section 5 discusses the ice abundances as well as the variations in ice mantle composition derived from fits to representative CO$_2$ ice profiles using current laboratory ice analogs. | We used new {\it Spitzer} data obtained with the IRS instrument to produce fully sampled maps of the distribution of CO$_2$ ices (15.2 $\mu$m), H$_2$O ices (6.02 $\mu$m) and total hydrogen nuclei, as inferred from the 9.7 $\mu$m silicate feature, toward the star-forming region Cepheus A East. We find that all solid state features peak at, and are distributed closely around, the spatial position of the deeply embedded protostar HW2 and coinciding with the NH$_3$(1,1) molecular clump CepA-1. Otherwise, the distributions show column densities lower by about a factor 2 than those measured over the CepA-1 region. The correlations we observe between the CO$_2$ ice, the H$_2$O ice column density distributions and that of total hydrogen imply that the solid state features predominantly arise from the same molecular clouds along the probed sight lines. We hence derived the CO$_2$ and water ice abundances with respect to the measured total hydrogen column density at each summed spatial position in the map. We find an average CO$_2$ ice abundance of (1.9 $\pm$ 0.4) $\times$ 10$^{-5}$ and an average H$_2$O ice abundance of (7.5 $\pm$ 1.7) $\times$ 10$^{-5}$ over the probed region. A comparison of the CO$_2$ and water ice column density distributions indicates that both ices build up onto dust grains, in concert, over the probed region. The fairly good correlation between the column densities of the two ices also indicates that blending of the water ice band by unidentified features is not significant in our data. Overall, we find that $N$(CO$_2$)$_{ice}=$ (0.22 $\pm$ 0.03) $\times$ $N$(H$_2$O)$_{ice}$, a fraction similar to that typically found in the intra-could medium and toward quiescent molecular clouds (17 \%), suggesting that grain surface chemistry is the most likely production mechanism of CO$_2$ ices toward this region. Best fits to the CO$_2$ ice bending mode using the current laboratory interstellar ice analogs database indicate that the ice mantles in Cepheus A East are composed of 2 ice mixtures, a low-temperature ($T_{lab}$=20K) polar ice mixture and a high-temperature ($T_{lab}$=119K) methanol-rich mixture. We find that while about 30\% of the probed ice mantles is at low temperature, variations of the cold-ice fraction also seem to exist over the probed region. The presence of the high-temperature methanol-rich ice mixture is indicative of significant thermal processing of a fraction of these ice mantles over Cepheus A East history. | 7 | 10 | 0710.0591 |
0710 | 0710.1903_arXiv.txt | We study the deflection of light in the background of a ``wiggly" cosmic string, and investigate whether it is possible to detect cosmic strings by means of weak gravitational lensing. For straight strings without small-scale structure there are no signals. In the case of strings with small-scale structure leading to a local gravitational attractive force towards the string, there is a small signal, namely a preferential elliptical distortion of the shape of background galaxies in the direction corresponding to the projection of the string onto the sky. The signal can be statistically distinguished from the signal produced by a linear distribution of black holes by employing an ellipticity axis distribution statistic. | There has been renewed interest in cosmic strings \cite{original} as a contributing factor to cosmological structure formation. This resurgence of interest is due firstly to the realization \cite{Rachel} that in many supersymmetric particle physics models, cosmic strings are formed after inflation, and thus contribute to but not completely replace inflationary perturbations as the seeds for structure formation. Secondly, it has recently been realized that models with cosmic superstrings \cite{Witten} may well be viable \cite{CMP}. They could, for example, be generated as the remnant of brane annihilation processes in brane inflation models \cite{Tye}, or they may play an important role in inflationary models in warped backgrounds \cite{stringinflation}. Cosmic superstrings may also be left behind after the initial Hagedorn phase in string gas cosmology \cite{BV}, where they would add an additional component to the spectrum of fluctuations produced by thermal string gas fluctuations \cite{NBV}. Unlike in the original cosmic string models of structure formation (see e.g. \cite{CSstructure} and \cite{CSreviews} for reviews), where it was assumed that the strings were the sole source of structure formation, in the context of the current models in which cosmic strings arise at the end of inflation, the strings contribute only a small fraction of the total power to the density fluctuations. The most stringent constraints on the fraction of the power which cosmic strings can contribute come from measurements of the angular power spectrum of coordinates cosmic microwave background (CMB) anisotropies. A large fraction $f$ of the power being due to strings is inconsistent with the observed acoustic oscillations in the angular power spectrum. The best current limits on $f$ are \cite{CMBlimit} $f < 10^{-1}$. Recent work \cite{ABB,Fraisse} points to the possibility that statistical analyses searching for the line discontinuities \cite{KS} in the microwave temperature maps produced by cosmic strings might lead to even tighter limits. Given their interest from the point of view of particle physics and superstring theory, it is of great interest to develop statistics to search for strings in observational data. There has been a substantial amount of work on identifying distinctive signals for strings in CMB temperature maps. In this paper we wish to take first steps at exploring another avenue - weak gravitational lensing. Weak gravitational lensing (see e.g. \cite{lensing} for a general review of the applications of gravitational lensing to cosmology) is emerging as a powerful tool in observational cosmology to search for the distribution of matter in the universe. One of the advantages of weak gravitational lensing over, for example, galaxy redshift surveys, is that light deflection depends on the total mass, not just the luminous mass. As a consequence, cosmic strings (which are also dark in the sense of not emitting light) also lead to gravitational lensing. In fact, due to conical structure of the metric of space-time in the plane perpendicular to a long string \cite{deficit}, the specific lensing pattern produced by a string can lead to interesting strong lensing patterns, e.g. double images (see \cite{Esther,recent} for claims to have detected such events, claims which were subsequently shown to be incorrect \cite{true}). In this short paper, we take a first step at searching for weak lensing signals from strings. Cosmic strings \cite{original} are one-dimensional topological defects which arise during phase transitions in the very early universe. Since they carry energy, they will lead to density fluctuations and CMB anisotropies. Causality implies that the network of strings which forms during the phase transition contains infinite strings. Once formed in the early universe, the network of strings will approach a ``scaling solution'' which is characterized by of the order one infinite string segment in each Hubble volume, and a distribution of cosmic string loops which are the remnants of the previous evolution. In particular, this implies that in any theory which admits cosmic strings, a network of strings will be present at the current time. The network will consist of a small number of strings crossing out entire Hubble volume (these are commonly called the ``infinite'' strings). The mean curvature radius $R_c$ of these strings will be comparable to the Hubble radius. Thus, for observations on smaller angular and distance scales, these strings can be approximated as straight. In addition, there is a distribution of string loops with radii $R$ between $R_c$ and a lower scale set by the strength of the gravitational radiation which the loops emit \cite{CSstructure}: \be G \mu R_c \, \le \, R \, \le \, R_c \, . \ee String loops with smaller radius live for less than one Hubble expansion time. In the above, $\mu$ is the mass per unit length of the string and $G$ is Newton's gravitational constant The distinctive signatures of strings in observational data are a consequence of the specific form of the metric produced by a cosmic string: space perpendicular to a cosmic string is a cone with deficit angle given by \cite{deficit} \be \alpha \, = \, 8 \pi G \mu \, . \ee In this paper we study the specific signature of this cosmic string metric on weak gravitational lensing. We consider a cosmic string lens and investigate the shape distortions this lens induces for a screen of background galaxies as source objects. To our knowledge, this is the first investigation of weak gravitational lensing from cosmic strings. Strong lensing signatures (lines of double images) have been studied extensity \cite{Vilenkin2,Hogan,Gott,Paczynski,Dyer,Vachaspati,Mack} theoretically, and looked for in observational data sets \cite{Hindmarsh,Shirasaki}. The outline of this paper is as follows: we first derive the deflection angle induced by a cosmic string with small scale structure. We then describe a gravitational lensing simulation based on this result, and propose a statistic which could be used to search for this type of effect in future weak lensing surveys. | In this work we have studied the gravitational lensing by a straight cosmic string containing small-scale structure which leads to a string tension which is less than the mass per unit length of the string, and thus induces a net gravitational force towards the string which test particles feel. Next, we studied potential weak lensing signatures of such wiggly strings. We found a shape distortion which is proportional to $G (\mu - T)$ but depends on the location of the source relative to the line of sight between the observer and the string. Only objects which are displaced in direction of the string relative to the line of sight give rise to a shape distortion. The shape distortion increases linearly as the distance in string direction increases. The specific distribution of the weak lensing distortion in the image plane in principle can be used to provide a signature for cosmic strings. However, in practise the signal appears to be too small to be useful. Even for optimal choices of parameters (galaxy redshifts comparable to those of galaxies in the largest current redshift surveys), distance from the line of sight comparable to the distance of the string from the observer, large amount of small scale structure and string tension close to the current observational bounds, the intrinsic ellipticity of the background galaxies needs to be very close to one (i.e. the shape needs to be very close to spherical) in order for the cosmic string signal to stand out. | 7 | 10 | 0710.1903 |
0710 | 0710.3906_arXiv.txt | We model the thermal effect of young stars on their surrounding environment in order to understand clustered star formation. We take radiative heating of dust, dust-gas collisional heating, cosmic-ray heating, and molecular cooling into account. Using Dusty, a spherical continuum radiative transfer code, we model the dust temperature distribution around young stellar objects with various luminosities and surrounding gas and dust density distributions. We have created a grid of dust temperature models, based on our modeling with Dusty, which we can use to calculate the dust temperature in a field of stars with various parameters. We then determine the gas temperature assuming energy balance. Our models can be used to make large-scale simulations of clustered star formation more realistic. | Most of the stars in our galaxy form in groups or clusters \citep{lada}. Therefore, in order to understand the star formation history, the shape of the mass function, and the formation of massive (M $\ga 5$ M$_{\odot}$) stars in our galaxy, the star formation process must be studied in its most common environment -- a cluster. As stars form from their initial reservoir of gas and dust, they interact with their environment and heat the surrounding material, thus affecting future star formation. One of the first effects a protostar has on its environment is radiative heating from the accretion luminosity and, subsequently, nuclear fusion. The radiation efficiently heats the dust, which in turn heats the gas through collisions. Young stars also affect their environment via strong winds and ionization, but this only occurs when they are very massive and have evolved past the very early stages of star formation. We assume that the massive stars in our sample are very young and are accreting at very high rates ($\dot{\rm{M}} \ga 1 \ee{-5} \msun / $yr). This high accretion rate allows the infalling mass to absorb all of the stellar UV photons \citep{Churchwell}. Many groups use large scale computer simulations to model clustered star formation. This is a complicated process requiring many assumptions in order to make the problem tractable. \cite{klessen} and Martel, Evans, \& Shapiro (2006) assume that the gas is isothermal. \citet{bbb} go beyond this assumption by using a barotropic equation of state. However, until recently, no one has included the effect of radiatively heating the dust and gas by the stars formed in the simulation. \citet{krumholz} have included an approximate radiative transfer method, which works well in optically thick regions. Their method assumes that the gas temperature is equal to the dust temperature throughout their simulation. This approximation is only valid at high densities when the dust and gas are collisionally coupled. Here we develop a method that explores the effect of radiative heating and the dust and gas energetics for a range of optical depths and densities. In our method we include various heating and cooling processes to calculate the dust and gas temperature. Stars can heat dust grains more effectively than the gas because dust grains have broad-band absorption properties. Although we will not be explicitly modeling the motion or energy density of dust grains, we assume the dust and gas are well-mixed and the dust grains transfer energy to gas particles through collisions using the energy transfer rate discussed in \citet{young}. The gas is heated by collisions with hot dust grains and cosmic rays. It can cool through CO and other molecular line emission. In this paper, we calculate the dust and gas temperature in a field of stars. The dust temperature around a single source is calculated using a look-up table which we develop here. With this look-up table and an approximation to the flux-temperature conversion, we calculate the dust temperature in the field. Our look-up table is needed since the calculation of a single dust temperature distribution can take longer than a minute on current desktops and would take a substantial fraction of a large-scale simulation's computations. Therefore we outline our method here which can be used to decrease the time spent on the calculation of the dust temperature in future studies of clustered star formation. With the calculated value of the dust temperature, we derive the gas temperature field for a distribution of stellar sources, as in a young stellar cluster. The effect that protostars have on heating their environment using a hydrodynamic and gravity simulation will be addressed in a future paper. In this paper, we first discuss the calculation of the dust temperature for single and multiple sources (\S \ref{sec:analy}), then we describe our gas temperature calculation (\S \ref{sec:gas}), and finally, we show some dust and gas temperature distributions when multiple sources are present (\S \ref{sec:multiple}). | We have presented a method for calculating the dust and gas temperature between stellar sources. The analytic method that we investigated for calculating the dust temperature was not accurate enough. Instead, our chosen method of calculating the dust temperature uses a simple radiative transfer code which we use to create a look-up table. Once we have derived the dust temperature, we are able to calculate the gas temperature by balancing various energy processes. We include dust-gas collisional heating, molecular cooling, and cosmic-ray heating. When we have balanced the energies, we are able to derive the gas temperature. Other methods which set the gas temperature and dust temperature equal assume the gas and dust are opaque everywhere, which is not always true. In Figure \ref{fig:perdiff}, we show the percentage difference between the gas and dust temperature, as well as the density, for the distribution of sources discussed in \S \ref{sec:three}. These two figures show that the largest percentage difference between the dust and gas temperature occurs where the density is the lowest. Therefore, at low densities ($n \la 10^{5}$cm$^{-3}$), assuming equal dust and gas temperatures is not appropriate. We plan to use the method discussed in this paper to model a region of clustered star formation with the three-dimensional hydrodynamics code discussed in \cite{martel}. Our method of calculating the gas and dust temperature distribution in a field of young stars will enable us and others to more accurately model clustered star formation observationally and in future simulations. | 7 | 10 | 0710.3906 |
0710 | 0710.3517.txt | Forty years after the discovery of rotation-powered pulsars, we still do not understand many aspects of their pulsed emission. In the last few years there have been some fundamental developments in acceleration and emission models. I will review both the basic physics of the models as well as the latest developments in understanding the high-energy emission of rotation-powered pulsars, with particular emphasis on the polar-cap and slot-gap models. Special and general relativistic effects play important roles in pulsar emission, from inertial frame-dragging near the stellar surface to aberration, time-of-flight and retardation of the magnetic field near the light cylinder. Understanding how these effects determine what we observe at different wavelengths is critical to unraveling the emission physics. I will discuss how current and future X-ray and gamma-ray detectors can test the predictions of these models. | \label{sec:intro} % Always give a unique label % and use \ref{<label>} for cross-references % and \cite{<label>} for bibliographic references % use \sectionmark{} % to alter or adjust the section heading in the running head Rotation-powered pulsars are fascinating astrophysical sources and excellent laboratories for study of fundamental physics of strong gravity, strong magnetic fields, high densities and relativity. The major advantage we have in studying pulsars is that we know they are rotating neutron stars and that they derive their power from rotational energy loss. The challenge is then to understand how they convert this source of power into the visible radiation. It is generally agreed that this occurs through acceleration of charged particles to extremely relativistic energies, using the rotating magnetic field as a unipolar inductor to create very high electric potentials. Beyond this fundamental, there is a large divergence of thought on what comes next: whether the acceleration occurs in the strong field near the neutron star surface or in the outer magnetosphere near the speed of light cylinder, or even beyond the light cylinder in the wind zone. The particle acceleration may well be occurring in all of these regions, either in the same pulsar or in pulsars of different ages. In recent years, there has been much activity both in new detections and in theoretical study of rotation-powered pulsars. Multibeam radio surveys at the Parkes Telescope \cite{Manchester2001,Edwards2001} have increased the population of known radio pulsars to more than 1700. In addition, extended radio observations of supernova remnants and unidentified $\gamma$-ray sources have discovered a number of young pulsars that are too radio-faint to be detected by surveys \cite{Camilo2004}. Although pulsed emission at other wavelengths has been detected from only a small fraction of these, this number is growing as well. At the present time, there are 7 pulsars with high-confidence detection of $\gamma$-ray pulsations \cite{Kanbach2006}, about 30 having X-ray pulsations \cite{Kaspi2006} and 10 with optical pulsations \cite{Mignami2004}. This paper will review both the fundamental physics and latest theoretical developments of acceleration and radiation in polar cap models. A complementary paper by Cheng \cite{Cheng2006} reviews acceleration and radiation in the outer gap model. | 7 | 10 | 0710.3517 |
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0710 | 0710.0029_arXiv.txt | Using the OVRO, Nobeyama, and IRAM mm-arrays, we searched for ``disk''-outflow systems in three high-mass (proto)star forming regions: \gs, \gtt, and \gte. These were selected from a sample of \amm\ cores (Codella, Testi \& Cesaroni) associated with OH and \wat\ maser emission (Foster \& Caswell) and with no or very faint continuum emission. Our imaging of molecular line (including rotational transitions of \mcn) and 3\,mm dust continuum emission revealed that these are compact ($\lesssim$ 0.05 -- 0.3 pc), massive ($\sim$ 100 -- 400 \Msun), and hot ($\sim$100 K) molecular cores (HMCs), that is likely sites of high-mass star formation prior to the appearance of ultracompact \hii\ regions. All three sources turn out to be associated with molecular outflows from \co\ and/or HCO$^+$ $J=$1--0 line imaging. In addition, velocity gradients of 10 -- 100 \kms\ pc$^{-1}$ in the innermost ($\lesssim$0.03 -- 0.13 pc), densest regions of the \gtt\ and \gte\ HMCs are identified along directions roughly perpendicular to the axes of the corresponding outflows. All the results suggest that these cores might be rotating about the outflow axis, although the contribution of rotation to gravitational equilibrium of the HMCs appears to be negligible. Our analysis indicates that the 3 HMCs are close to virial equilibrium due to turbulent pressure support. Comparison with other similar objects where rotating toroids have been identified so far shows that in our case rotation appears to be much less prominent; this can be explained by the combined effect of unfavorable projection, large distance, and limited angular resolution with the current interferometers. | \label{ss:intro} The role of disks in the formation process of low-mass stars (i.e. stars with masses $\lesssim 1$ \Msun) has been extensively studied in the last two decades through high angular resolution observations at various wavelengths. Images of such disks have been obtained in the optical (e.g. Burrows et al. 1996) and at mm-wavelengths (e.g. Simon et al. 2000). The latter have demonstrated that the majority of the disks undergo Keplerian rotation. These findings are consistent with the fact that low-mass stars form through accretion while (partial) conservation of angular momentum during the dynamical collapse produces a flattened and rotating structure at the center of the core.\par What about high-mass ($M_\ast \gtrsim 8$ \Msun) stars? In this case, formation through accretion faces the problem that stars more massive than $\sim$8~$M_\odot$ reach the zero age main sequence still deeply embedded in their parental cores (Palla \& Stahler 1993). At this point radiation pressure from the newly formed early-type star can halt the infall, thus preventing further growth of the stellar mass. Various solutions have been proposed to solve this problem: (i)~massive stars might form through merging of lower mass stars (Bonnell, Bate \& Zinnecker, H. 1998; Bonnell \& Bate 2002; Bally \& Zinnecker 2005); (ii)~sufficiently large accretion rates could allow the ram pressure of the infalling material to overcome the radiation pressure form the star (Behrend \& Maeder 2001; McKee \& Tan 2003; Bonnell, Vine, \& Bate 2004) (iii)~non-spherical accretion could weaken the effect of radiation pressure by allowing part of the photons to escape through evacuated regions along the outflow axis and, at the same time, enhancing the ram pressure of the accreting material by focusing it through the disk plane (Yorke \& Sonnhalter 2002; Krumholz, McKee, \& Klein 2005).\par An important test to discriminate between the different hypotheses is the presence of rotating, circumstellar disks, which would lend support to the third scenario depicted above. This can be determined by inspecting the velocity field of the innermost parts of molecular cores where young massive (proto)stars are believed to form. One possibility is to look for velocity gradients perpendicular to the direction of the larger scale outflows associated with such cores; this would suggest that one is observing rotation about the outflow axis. This technique has been adopted by us and other authors successfully inferring the existence of rotation in the gas enshrouding high-mass young stellar objects (YSOs) and leading to the discovery of circumstellar disk-like structure or ``toroids'' in a limited number of cases (see Cesaroni et al. 2007 for a review on this topic).\par The goal of the present study was to establish whether the presence of rotation is common in high-mass star forming regions by observing three more objects, expected to be sites of deeply embedded OB (proto)stars. The ideal target for this type of studies are hot molecular cores (HMCs), which are believed to be the cradles of massive stars (see e.g. Kurtz et al. 2000; Cesaroni et al. 2007). Beside other studies, the one by Codella, Testi \& Cesaroni (1997, hereafter CTC97), has proved successful in identifying the birthplaces of young massive stars as dense \amm\ cores associated with OH and H$_2$O maser emission. One of these HMCs (G\,24.78+0.08) has been the subject of a series of articles by us (Furuya et al. 2002; Cesaroni et al. 2003; Beltr\'an et al. 2004, 2005), which have shown that this is a unique object, characterized by the simultaneous presence of rotation, outflow, and infall towards a hypercompact \hii\ region ionized by an O9.5 star (Beltr\'an et al. 2006). With this in mind, we completed the study of the sources in CTC97 by observing three more objects in the same tracers used to investigate G\,24.78+0.08. The selected targets for this study are G\,16.59$-$0.05 at $d =$ 4.7 kpc, G\,28.87$+$0.07 at 7.4 kpc, and G\,23.01$-$0.41 at 10.7 kpc. The fact that in all three cases no or only faint free-free continuum emission has been detected, notwithstanding the large luminosities of the sources (CTC97), suggests that the embedded YSOs might be in an even earlier evolutionary phase than those in G\,24.78+0.08 (Furuya et al. 2002; Beltr\'an et al. 2004). | Using the OVRO, Nobeyama, and IRAM-PdB mm-interferometers, we carried out intensive search for rotating toroids towards the massive YSOs in \gs, \gte, and \gtt\ which exhibit no or faint free-free emission (CTC97). Our observations revealed that these objects are embedded in HMCs with masses of 95--380 \Msun\ and temperatures of 93--130 K, making the cores typical site of high-mass (proto)star formation. All the 3 objects harbored in the HMCs are driving powerful (\Fco $\simeq$ $10^{-3}-10^{-2}$ \Msun \kms\ yr$^{-1}$) CO outflows. However, the nature of the outflows in \gte\ and \gtt\ is unclear; the origin of high velocity wing emission may be attributed to either single or double outflow(s). Such ambiguity made the interpretation of velocity gradients, identified through \mcn\ $K$-ladder line analysis, existing in the innermost densest part of the \gte\ and \gtt\ HMCs fairly difficult. The velocity gradients are almost perpendicular to their molecular outflow axes, suggesting the presence of rotating, flattened structures. However, the corresponding dynamical masses are an order of magnitude smaller than the masses derived from 3\,mm dust continuum emission, which indicates turbulent pressure as the dominant support of the HMCs. No conclusion could be reached for the third source, \gs, as the putative rotation axis appears to lie close to the line-of-sight, thus making the detection of the rotation velocity very difficult for projection effects. Further higher resolution imaging will allow us to establish the presence of rotation on a more solid ground. | 7 | 10 | 0710.0029 |
0710 | 0710.3670_arXiv.txt | We discuss a general fourth-order theory of gravity on the brane. In general, the formulation of the junction conditions (except for Euler characteristics such as Gauss-Bonnet term) leads to the higher powers of the delta function and requires regularization. We suggest the way to avoid such a problem by imposing the metric and its first derivative to be regular at the brane, while the second derivative to have a kink, the third derivative of the metric to have a step function discontinuity, and no sooner as the fourth derivative of the metric to give the delta function contribution to the field equations. Alternatively, we discuss the reduction of the fourth-order gravity to the second-order theory by introducing an extra tensor field. We formulate the appropriate junction conditions on the brane. We prove the equivalence of both theories. In particular, we prove the equivalence of the junction conditions with different assumptions related to the continuity of the metric along the brane. | \label{sect1} \setcounter{equation}{0} Brane universes have made great popularity during the last years \cite{RS,brane}. However, it is remarkable that so far only the standard Einstein gravity, Gauss-Bonnet gravity \cite{deruelle00,charmousis,davis,jim,lidsey,maeda,apostopoulos} and, in general, Euler density gravity \cite{lovelock} on the brane have been considered in the literature \cite{meissner01}. These can be expressed by the general action \footnote{We use convention (-++...+) for the metric following Ref. \cite{he}.} \bea \label{euler} S = \int_M d^D x \sqrt{-g} \sum_n \kappa_n I^{(n)} + S_{brane} + S_{m}~, \eea where $I^{(n)}$ is the Euler density of the n-th order, $\kappa_n$ is an appropriate constant of the n-th order, $M$ is a $D$-dimensional manifold, $S_{brane}$ is the brane action and $S_m$ is the matter action. The lowest order Euler densities are: the cosmological constant $I^{(0)} = 1$, the Ricci scalar $I^{(1)} = R$, and the Gauss-Bonnet density $I^{(2)} = R_{GB} = R_{abcd}R^{abcd} - 4 R_{ab}R^{ab} +R^2$ with appropriate constants $\kappa_0 = -2\Lambda(2\kappa^2)^{-1} = -2\Lambda/16 \pi G$, $\kappa_1 = (2\kappa^2)^{-1}$, $\kappa_2=\alpha(2\kappa^2)^{-1}$, $\alpha=$ const. etc., $a,b,c=0,1,\ldots, D-3, D-2, D$ \cite{davis}. In fact, in a general class of brane models based on an arbitrary combination of the higher-order curvature terms $f(R_{abcd}R^{abcd},R_{ab}R^{ab},R)$ the field equations are fourth-order. Because of that they are plagued by the higher power terms of the second derivative of the warp factor function $\sigma(y) = \mid y \mid$. This leads to a production of the higher powers of the delta function $\delta(y)$ which can make the field equations ambiguous. Among the general class, the models based on the Euler densities are unique in the sense that the higher powers of the second derivative of the warp factor $\partial^2 \sigma(y)/\partial y^2$ exactly cancel in the field equations \cite{meissner01}. One then is easily able to formulate appropriate junction conditions given first by Israel \cite{israel66,visser}. The main objective of our paper will be the study of a general {\it fourth-order theory of gravity on the brane} \cite{clifton} \bea \label{XYZ} S &=& \chi^{-1} \int_{M} d^{D}x \sqrt{-g} f(X,Y,Z) + S_{brane} + S_{m}~, \eea where $X=R$, $Y=R_{ab}R^{ab}$, $R_{abcd}R^{abcd}$ are curvature invariants, and $\chi$ is a constant. It includes the Euler density theories with the first Euler density being just $f(X,Y,Z) = \chi \kappa_1 X = \chi \kappa_1 R $ and the second Euler density being the Gauss-Bonnet term given by $f(X,Y,Z) = \chi \kappa_2 (Z - 4Y +X^2)$ etc. Up to our knowledge, the only non-Eulerian density cases were studied in Refs. \cite{branef(R)} and \cite{braneR2}. In Ref. \cite{branef(R)} the fourth-order theory $f(X,Y,Z) = f(X) = f(R)$ was first reduced to the second-order theory, and then transformed into the Einstein theory. The junction conditions were then obtained, and they were obviously free from the problem of the powers of $\delta-$function contribution. On the other hand, in Ref. \cite{braneR2} the theories with the linear combination of the form $f(X,Y,Z) = aX^2 + bY + cZ$ ($a,b,c=$ const.) were considered and the junction conditions were obtained by the application of the appropriate Gibbons-Hawking boundary terms, again after transforming this theory to an equivalent second-order theory. It is important to emphasize that the theories based on the functions of the Euler densities such as $f(I^{(n)})$ are fourth-order. Among them the most popular are $f(I^{(1)}) = f(R)$ - the theories of the function of the first Euler density \cite{f(R)}. In fact, the theories which are based on the function of the second Euler density $f(I^{(2)})$ have also gained some interest recently \cite{f(RGB)}, but they have not been studied on the brane yet. Our paper is organized as follows. In Section II we discuss the main obstacle to formulate junction conditions for the fourth-order braneworld in a standard way which has been performed in the case of Euler densities. In Section III we make a proposal to formulate these junction conditions by imposing more regularity onto the metric tensor. Since it does not necessarily satisfy everybody's taste we present in Section IV an alternative approach. In this approach we transform our general fourth-order theory into a second-order theory by applying generalized Lagrange-multiplier approach \cite{f(R),f(RGB),kijowski}. This method was successful in obtaining the junction conditions in $f(R)$ theory \cite{branef(R)} and in $f = aX^2 + bY + cZ$ theory \cite{braneR2}. In the Section V we formulate the junction conditions for the equivalent second-order theory. Finally, in Section VI we give our conclusions. | \label{sum} \setcounter{equation}{0} In this paper we have considered the junction conditions for the fourth-order brane world given by the action which is a general function of the curvature invariants $R,R_{ab}R^{ab}$, and $R_{abcd}R^{abcd}$. We imposed the regularity conditions on the metric tensor which was taken to be of the class $ C^2$ functions of the coordinates. Explicitly, it means that the metric and its first derivative are regular at the brane, while the second derivative has a kink, the third derivative of the metric has a step function discontinuity, and the fourth derivative of the metric gives the delta function contribution to the field equations. In terms of the seminal notation given first by Israel \cite{israel66}, these conditions describe the singular hypersurfaces of order three. The junction conditions which we obtained are quite generic and they may be applied to some special cosmological framework of interest. As an alternative which allows less restrictive regularity conditions, we considered the reduction of the fourth-order theory to a second-order theory by applying an extra tensor field. We then formulated the junction conditions within such a theory and showed that they were equivalent to the previously obtained fourth-order theory junction conditions. In the previously considered cases in the literature mainly the Eulerian density brane worlds were studied \cite{deruelle00,charmousis,davis,jim,lidsey,maeda,apostopoulos,meissner01}. The only non-Eulerian density cases were investigated in Refs. \cite{branef(R)} and \cite{braneR2}, where the theories with at most a linear combination the curvature invariants of the form $f(X,Y,Z) = aX^2 + bY + cZ$ ($a,b,c=$ const.) were considered. In these references the junction conditions were obtained after transforming such theories into the equivalent second-order theories. In fact, we made one step further, by suggesting junction conditions for the brane world which allows a general function of curvature invariants $f(X,Y,Z)$. Also, in one of our approaches to the problem of junction conditions we do not reduce the fourth-order action into the second-order action, but suggest junction conditions to work in the fourth-order theory, though in the case of the so-called singular hypersurfaces of order three only \cite{israel66}. We hope that our calculations will allow us to study some cosmological applications of the general fourth-order gravity on the brane (\ref{XYZ}) in the following papers. \vspace{0.3cm} | 7 | 10 | 0710.3670 |
0710 | 0710.1958_arXiv.txt | We explore the properties of dark energy from recent observational data, including the Gold Sne Ia, the baryonic acoustic oscillation peak from SDSS, the CMB shift parameter from WMAP3, the X-ray gas mass fraction in cluster and the Hubble parameter versus redshift. The $\Lambda CDM$ model with curvature and two parameterized dark energy models are studied. For the $\Lambda CDM$ model, we find that the flat universe is consistent with observations at the $1\sigma$ confidence level and a closed universe is slightly favored by these data. For two parameterized dark energy models, with the prior given on the present matter density, $\Omega_{m0}$, with $\Omega_{m0}=0.24$, $\Omega_{m0}=0.28$ and $\Omega_{m0}=0.32$, our result seems to suggest that the trend of $\Omega_{m0}$ dependence for an evolving dark energy from a combination of the observational data sets is model-dependent. | The present cosmic accelerating expansion has been confirmed by various observations, including the Type Ia Supernovae (Sne Ia)~\cite{Perlmutter1999, Riess1998, Riess2004, Riess2006, Astier2006, Wood2007}, CMB \cite{Balbi2000, de2000, Jaffe2001,Spergel2003, Spergel2006} and large scale structure (LSS) \cite{Peacock2001, Eisenstein2005}, etc. In order to explain this observed phenomenon, it is usually assumed that there exists, in the universe, an exotic energy component with negative pressure, named dark energy (see \cite{Padmanabhan2006, Copeland2006, Sahni2006, Perivolaropoulos2006} for recent reviews), which presumably began to dominate the evolution of the universe only recently. The simplest candidate of dark energy is the cosmological constant $\Lambda$~\cite{Weinberg1989, Sahni2000, peebles2003, Padmanabhan2003}. It fits the observational data very well, but at the same time, it also encounters two problems, i.e., the cosmological constant problem (why is the inferred value of cosmological constant so tiny ($120$ orders of magnitude lower) compared to the typical vacuum energy values predicted by particle physics?) and the coincidence problem (why is its energy density comparable to the matter density right now?). Therefore, some dynamical scalar fields, such as quintessence~\cite{Wetterich1988, Ratra1988, Caldwell1998}, phantom~\cite{Cald} and quintom~\cite{Quintom}, etc, are suggested as alternative candidates of dark energy. One of the features of these scalar field models is that their equations of state parameter, $w$, which embodies both gravitational and evolutionary properties of dark energy, is evolving with the cosmic expansion. On the other hand, the growing number of dark energy models has prompted people to adopt a complementary approach, which assumes an arbitrary parametrization for the equation of state $w(z)$ in a model-independent way and aims to reconstruct the properties of dark energy directly from observations. Currently, there are many model independent parameterizations (see for example, ~\cite{line1,line2,Jassal2005,Sahni2003, Sahni2004, Wetterich2004}). In general, using these parameterizations and the observational data, one can determine the present value of $w$ and whether it evolves as the universe expands, in particular, whether the phantom divide line (PDL) is crossed. In this regard, Nesseris and Perivolaropoulos \cite{Nesseris2006} has used the Chevallier-Polarski-Linder parametrization $ w(z) = w_0 + w_1 z/(1 + z)$~\cite{line2} to explore the properties of dark energy with some observational data (including new Gold Sne Ia, SNLS Sne Ia, CMB, BAO, the cluster baryon gas mass fraction(CBF) and 2dF galaxy redshift survey(2dFGRS) ) and found that the Gold data set mildly favors dynamically evolving dark energy with the crossing of the PDL while the SNLS does not, and the combination of CMB+BAO+CBF+2dFGRS mildly favors the crossing of PDL only for low values of $\Omega_{m0}$ ( $\Omega_{m0}\le 0.25$) prior considered and with a higher prior matter density the evolving features of dark energy becomes weaker and weaker. Similar trend of $\Omega_{m0}$ dependence was found using the model~\cite{Alam2006}, $w(z)=\frac{1+z}{3}\frac{A_1+2A_2(1+z)}{\Omega_{DE}}-1$, with the CMB and BAO. However, constraints from a combination of the supernovae and other observational data has not been analyzed in Ref.~\cite{Nesseris2006}, and although that of the Sne and CMB+BAO was examined in Ref.~\cite{Alam2006}, but the marginalization was considered only for $\Omega_{m0}=0.28\pm 0.03$ prior. Therefore, it remains interesting to see what happens to the conclusions reached in Refs.~\cite{Nesseris2006,Alam2006}, when the combination of all observational data is analyzed for different $\Omega_{m0}$ prior considered. The present paper aims to fill the gap. We discuss the constraints from the combination of different observational datasets. Besides the data sets of Sne Ia, BAO and CMB, in our analysis we add the datasets of the X-ray gas mass fraction in cluster and the Hubble parameter versus redshift. Firstly the $\Lambda CDM$ model with curvature is discussed. Then, two parameterized dark energy models: $ w(z) = w_0 + w_1 z/(1 + z)$ and $w(z)=\frac{1+z}{3}\frac{A_1+2A_2(1+z)}{\Omega_{DE}}-1$, are studied to see if the properties of dark energy thus reconstructed are model-independent. | In this paper, we have reconstructed the properties of dark energy from recent observational data, including the Gold Sne Ia, the baryonic acoustic oscillation peak from SDSS, the CMB shift parameter, the X-ray gas mass fraction in clusters and the Hubble parameter data. The $\Lambda CDM$ model with curvature and two parameterized dark energy models are discussed. We find that a spatially flat universe is allowed by these data sets at the $68\%$ confidence level, and a closed universe is slightly favored by the observations. For two parameterized dark energy models, we give the priors on $\Omega_{m0}$ with $\Omega_{m0}=0.24$, $\Omega_{m0}=0.28$ and $\Omega_{m0}=0.32$. For the spatially flat case, the constraints on model parameters and the evolutions of $w(z)$ and $q(z)$ are studied. The Gold + CMB +BAO give the strong constraints on model parameters. If Mod1 parametrization is used, the best fit curves in Fig.~3 show that the combination of the data sets considered in this paper favors an evolving dark energy, a crossing of the phantom divide line in the near past, and the present value of $w$ being very likely less than $-1$. Remarkably, these conclusions are almost insensitive to the chosen value of matter density, in a sharp contrast to those obtained in~\cite{Nesseris2006} where the Sne data are not combined with other observational ones. However, the best fit curves in Fig.~3 indicate that the properties of dark energy reconstructed using Mod2 parametrization depend on the chosen value of matter density. For $\Omega_{m0}=0.24$, a very mildly evolving dark energy is obtained, but with the increase of $\Omega_{m0}$ prior considered, the evolving feature of dark energy becomes more evident. This trend is just the opposite to that found in Ref. \cite{Alam2006} for just BAO+CMB data. Therefore, our result seems to suggest that the trend of $\Omega_{m0}$ dependence for an evolving dark energy is model-dependent. It should be noted, however, that at the $2\sigma$ confidence level, the cosmological constant are allowed for both Mod1 and Mod2. | 7 | 10 | 0710.1958 |
0710 | 0710.0115_arXiv.txt | Null Energy Condition (NEC) requires the equation of state (EoS) of the universe $w_u$ satisfy $w_u\geq-1$, which implies, for instance in a universe with matter and dark energy dominating $w_u=w_m\Omega_m+w_{de}\Omega_{de}=w_{de}\Omega_{de}\geq-1$. In this paper we study constraints on the dark energy models from the requirement of the NEC. We will show that with $\Omega_{de}\sim0.7$, $w_{de}<-1$ at present epoch is possible. However, NEC excludes the possibility of $w_{de}<-1$ forever as happened in the Phantom model, but if $w_{de}<-1$ stays for a short period of time as predicted in the Quintom theory NEC can be satisfied. We take three examples of Quintom models of dark energy, namely the phenomenological EoS, the two-scalar-field model and the single scalar model with a modified Dirac-Born-Infeld (DBI) lagrangian to show how this happens. | It is well known that energy conditions play an important role in classical theory of general relativity and thermodynamics\cite{Hawking:1973uf}. In classical general relativity it is usually convenient and efficient to restrict a physical system to satisfy one or some of energy conditions for study, for example, in the proof of Hawking-Penrose singularity theorem\cite{Penrose:1964wq,Hawking:1969sw}, the positive mass theorem\cite{Schon:1981vd} and so on, while in thermodynamics energy conditions are the bases for obtaining entropy bounds\cite{Bousso:1999xy,Flanagan:1999jp}. Among those energy conditions, the null energy condition is the weakest one which states that for any null vector $n^\mu$ the stress energy tensor $T_{\mu\nu}$ should satisfy the relation \begin{eqnarray} T_{\mu\nu} n^\mu n^\nu \geq 0~. \end{eqnarray} In general, the violation of NEC leads to the breakdown of causality in general relativity and the violation of the second law of thermodynamics\cite{ArkaniHamed:2007ky}. These pathologies require that the total stress tensor in a physical spacetime manifold should obey the NEC. In the framework of the standard 4-dimensional Friedmann-Robertson-Walker (FRW) cosmology the NEC implies $\rho+p\geq 0$, which in turn gives rise to a constraint on the equation of state of the universe (EoS) $w_u$ defined as the ratio of pressure to energy density, $w_{u}\geq -1$. In this paper we study the constraints on the dark energy models from the requirement of $w_{u}\geq -1$. In the early Universe with radiation dominant the EoS of the universe $w_{u}$ is approximately equal to $\frac{1}{3}$ and in the matter dominant period $w_{u}$ is nearly zero, so NEC is satisfied well. However when the dark energy component is not negligible we have \begin{eqnarray}\label{NECd} w_u=w_m\Omega_m+w_{de}\Omega_{de}\geq-1~, \end{eqnarray} where the subscripts `$m$' and `$de$' stand for matter and dark energy, respectively. With $w_m = 0$, inequality (\ref{NECd}) becomes \begin{eqnarray}\label{NECde} w_{de}\Omega_{de}\geq-1~. \end{eqnarray} From the inequality above, we can see that models of dark energy with $w_{de}\geq-1$ such as the Cosmological Constant and the Quintessence satisfy the NEC, but the models with $w_{de}<-1$ predicted for instance by the Phantom theory where the kinetic term of the scalar field has a wrong sign does not. Interestingly we can see that NEC might be satisfied in models if $w_{de}<-1$ stays for a short period of time during the evolution of the universe. In this paper we will show this happens in the Quintom models of dark energy. The paper is organized as follows: in section II we will present three examples of the Quintom models to show how the NEC is satisfied and the section III is the summary of the paper. | In this paper we have studied the implications of NEC in the models of dark energy. We show that NEC excludes the models with $w_{de}<-1$ forever as predicted by the Phantom dark energy, however allows the possibility of having $w_{de}<-1$ for a short period of time as it happens in the Quintom models. We have shown explicitly in this paper three examples of Quintom models where NEC is satisfied. | 7 | 10 | 0710.0115 |
0710 | 0710.5440_arXiv.txt | We describe numerical tools for the stability analysis of extrasolar planetary systems. In particular, we consider the relative Poincar\'e variables and symplectic integration of the equations of motion. We apply the tangent map to derive a numerically efficient algorithm of the fast indicator MEGNO (a measure of the maximal Lyapunov exponent) that helps to distinguish chaotic and regular configurations. The results concerning the three-planet extrasolar system HD~37124 are presented and discussed. The best fit solutions found in earlier works are studied more closely. The system involves Jovian planets with similar masses. The orbits have moderate eccentricities, nevertheless the best fit solutions are found in dynamically active region of the phase space. The long term stability of the system is determined by a net of low-order two-body and three-body mean motion resonances. In particular, the three-body resonances may induce strong chaos that leads to self-destruction of the system after Myrs of apparently stable and bounded evolution. In such a case, numerically efficient dynamical maps are useful to resolve the fine structure of the phase space and to identify the sources of unstable behavior. | {Understanding } the extrasolar planetary systems has became a major challenge for contemporary astronomy. One of the most difficult problems in this field concerns the orbital stability of such systems. Usually, the investigations of long-term evolution are the domain of direct, numerical integrations. The stability of extrasolar systems is often understood in terms of the Lagrange definition implying that orbits remain well bounded over an arbitrarily long time. Other definitions may be formulated as well, like the astronomical stability \citep{Lissauer1999} requiring that the system persists over a very long, Gyr time-scale, or Hill stability \citep{Szebehely1984} that requires the constant ordering of the planets. In our studies, we prefer a more formal and stringent approach related to the fundamental Kolmogorov-Arnold-Theorem (KAM), see \cite{Arnold1978}. Planetary systems, involving a dominant mass of the parent star and significantly smaller planetary masses, are well modeled by close-to-integrable, Hamiltonian dynamical systems. It is well known, that their evolution may be quasi-periodic (with a discrete number of fundamental frequencies, forever stable), periodic (or resonant; stable or unstable) or chaotic (with a continuous spectrum of frequencies, and unstable). In the last case, initially close phase trajectories diverge exponentially, i.e., their Maximum Lyapunov Characteristic Exponent (MLCE, denoted also with $\sigma$) is positive. In general, the distinction between regular and chaotic trajectories is a very difficult task that may be resolved only with numerical methods relying on efficient and accurate integrators of the equations of motion. The main motivation of this paper is to describe numerical tools that are useful for studies of the dynamical stability and to apply them to the HD~37124 system \citep{Vogt2005}. We recall the fundamentals of relative canonical Poincar\'e variables as -- in our opinion -- one of the best frameworks for symplectic integrators. These {canonical} variables are well suited for the construction of a \citet{LasRob} composition method that improves a classical Wisdom-Holman (W-H) algorithm \citep{WH:91}. We supplement the integrator with a propagator of the associated symplectic tangent map that approximates the solution of variational equations \citep{MikIn:99}. Finally, we compare two fast indicators that reveal the character of phase trajectories. The first one is a relatively simple method for resolving fundamental frequencies and spectral properties of a close-to-integrable Hamiltonian system -- a so called Spectral Number (SN), invented by \citet{Michtchenko2001}. The second indicator belongs to the realm of the Lyapunov exponent based algorithms; we chose the numerical tool developed by \cite{Cincotta2000,Cincotta2003} under the name of MEGNO. In this work, we refine the algorithm of MEGNO that makes explicit use of the symplectic tangent map \citep{Gozdziewski2003b}. As a non-trivial application of the presented numerical tools, we consider the 3-planet system hosted by the HD~37124~star \citep{Vogt2005}. It has been discovered by the radial velocity (RV) technique. The recent model of the RV observations of HD~37124 predicts three equal Jovian type planets with masses $\sim 0.6$~m$_{\idm{J}}$ in orbits with moderate eccentricities. In such a case, the application of symplectic integrators without regularization is particularly advantageous thanks to the numerical efficiency (long time-steps) and accuracy (the total energy does not have a secular error and the angular momentum integral is conserved). The number of multi-planet systems resembling the architecture of the Solar system increases\footnote{For a recent statistics of the discoveries, see Jean Schneider's Extrasolar Planets Encyclopedia, http://exoplanets.eu.}. Hence, our approach may be useful in other cases. | The use of Poincar\'e variables in the studies of the dynamics of close-to integrable planetary systems offers many advantages. The variables are canonical and offer a simple form of a reduced Hamiltonian. The Hamiltonian can be split into a sum of two separately integrable parts: the Keplerian term and a small perturbation. As such, it can serve to construct a symplectic integrator based on any modern composition method, including the recent ones invented by \citet{LasRob}. The tangent map computed with the same integration scheme provides an efficient way of computing the estimate of maximal Lyapunov exponent in terms of relatively recent fast indicator MEGNO. The method proves to be much more efficient than general purpose integrators (like the Bulirsh-Stoer-Gragg method). Besides, it provides the conservation of the integrals of energy and the angular momentum that is crucial for resolving the fine structure of the phase space. From the practical point of view, the symplectic algorithms are relatively simple for numerical implementation. Using the numerical tools, we investigate the long term stability of extrasolar planetary system hosted by HD~37124. The orbital parameters in the set of our best, self-consistent Newtonian fits \citep{Gozdziewski2006a} are in accord with the discovery paper \citep{Vogt2005}. Nevertheless, the observational window of the system is still narrow and the derivation of the model consistent with observations is difficult and, in fact, uncertain. The dynamical maps reveal that the relevant region of the phase space, in the neighborhood of the mathematically best fit, is a strongly chaotic and unstable zone. The fitting algorithm (GAMP) that relies on eliminating strongly unstable fits founds solutions with a similar quality [in terms of $\Chi$] that yields the formal solution. Moreover, they are shifted towards larger semi-major axes and much smaller eccentricities of the outermost planet. The orbital evolution of two outer planets is confined to a zone spanned by a number of low-order two-body and three-body MMRs. In particular, the three-body MMRs may induce very unstable behaviors that manifest themselves after many Myrs of an apparently stable and bounded evolution. To deal with such a problem, the stability of the best fits should be examined over a time-scale that is much longer than the one required when only the two-body MMRs are considered. In accord with the dynamical maps, the stable fits to the RV of HD~37124 should have small eccentricity of the outermost planet~d, not larger than 0.2-0.3. Moreover, the stable configurations of the HD~37124 system are puzzling. The best-fit mathematical three-planet model is surprisingly distant, in the phase space of initial conditions, from the zone of stable solutions consistent with the RV. It remains possible that other bodies are present in the system and the three-planet model is not adequate to explain the RV variability, in spite that it provides apparently perfect fits. Yet, isolated initial conditions or even sets of best-fit solutions do not provide a complete answer on the system configuration. Then the fast indicator approach is essential and helpful to resolve the dynamical structure of the phase space. The results of our experiments confirm and warn that all numerical methods should be applied with great care. All symplectic methods are constant step integrators. In that case one should be cautious about the possibility of generating spurious resonance webs. A proper way to avoid them is to repeat computations with a different integration step in order to detect step-dependent patterns. | 7 | 10 | 0710.5440 |
0710 | 0710.0579_arXiv.txt | {} {We aim to provide observational constraints on diffusion models that predict peculiar chemical abundances in the atmospheres of Am stars. We also intend to check if chemical peculiarities and slow rotation can be explained by the presence of a weak magnetic field.} {We have obtained high resolution, high signal-to-noise ratio spectra of eight previously-classified Am stars, two normal A-type stars and one Blue Straggler, considered to be members of the Praesepe cluster. For all of these stars we have determined fundamental parameters and photospheric abundances for a large number of chemical elements, with a higher precision than was ever obtained before for this cluster. For seven of these stars we also obtained spectra in circular polarization and applied the LSD technique to constrain the longitudinal magnetic field.} {No magnetic field was detected in any of the analysed stars. HD~73666, a Blue Straggler previously considered as an Ap\,(Si) star, turns out to have the abundances of a normal A-type star. Am classification is not confirmed for HD~72942. For HD~73709 we have also calculated synthetic $\Delta a$ photometry that is in good agreement with the observations. There is a generally good agreement between abundance predictions of diffusion models and values that we have obtained for the remaining Am stars. However, the observed Na and S abundances deviate from the predictions by 0.6 dex and $\geq$0.25 dex respectively. Li appears to be overabundant in three stars of our sample.} {} | Main sequence A-type stars present spectral peculiarities, usually interpreted as due to peculiar photospheric abundances and abundance distributions which are believed to be produced by the interaction of a large variety of physical processes (e.g. diffusion, magnetic field, pulsation and various kinds of mixing processes). An interesting problem that has yet to be addressed is how these peculiarities change during main sequence evolution. The chemical composition of field A-type stars have been studied by several authors, e.g. \citet{hillland1993}, \citet{adelman2000}. However, it is not straightforward to use the results of these investigations to study how photospheric chemistry evolves during a star's main sequence life. First, the original composition of the cloud from which stars were born is not known and is likely somewhat different for each field star. It is therefore not possible to discriminate between evolutionary effects and differences due to original chemical composition. Secondly, it is difficult to estimate the age of field stars with the precision necessary for such evolutionary studies \citep[for a discussion of this problem see][]{stefano2006}. From this point of view, A-type stars belonging to open clusters are much more interesting objects. Compared to field stars, A-type stars in open clusters have three very interesting properties: \begin{list}{-}{} \item they were all presumably born from the interstellar gas with an approximately uniform composition; \item they all have approximately the same age (to within a few Myr); \item their age can be determined much more precisely than for field stars. \end{list} Few abundance analyses of A-type stars in open clusters have been carried out. Those that have been published have usually focused on a limited numbers of stars. \citet{monier1999} have determined the abundances of eleven chemical elements for a large sample of stars regularly distributed in spectral type along the main sequence in order to sample the expected masses uniformly. All these stars were analysed in a uniform manner using spectrum synthesis. \citet{ch2006} have performed a detailed abundance analysis for five A-type stars of the young open cluster IC~2391. \citet{folsom2007} have performed a detailed abundance analysis for four Ap/Bp stars and one normal late B-type star of the open cluster NGC~6475. A goal of this programme is to determine photospheric abundance patterns in A-type star members of clusters of different ages. This is crucial in order to: \textit{i)} investigate the chemical differences between normal and peculiar stars inside the same cluster, \textit{ii)} study the evolution with time of abundance peculiarities by studying clusters of various ages, \textit{iii)} set constraints on the hydrodynamical processes occurring at the base of the convection zone in the non magnetic stars and \textit{iv)} study the effects of diffusion in the presence of a magnetic field for the magnetic (Ap) stars in the cluster. The abundance analysis will be performed in an homogeneous way applying a method described in this first work. Praesepe (NGC~2632), a nearby intermediate-age open cluster \citep[$\log t = 8.85\pm 0.15$,][]{gonzalez}, is an especially interesting target because it includes a large number of A-type stars, among which are many Am stars. Furthermore, because the cluster is relatively close to the sun \citep[d~=~180~$\pm$~10 pc,][]{robichon1999}, many of the member A-type stars are bright enough to allow us to obtain high resolution spectra with intermediate class telescopes. We dedicate this first paper to the Am stars of the Praesepe cluster, searching for magnetic fields in these objects and discussing the differences between "normal"\footnote{We consider as "normal" A-type stars all the A-type stars that are classified neither as Am nor Ap} A-type stars and Am stars in the cluster. We also compare our results with previous works and with theoretical chemical evolution models. In particular we take into account diffusion models by \citet{richer2000}. We want to provide observational constraints to the theory of the evolution of the abundances in normal and chemically peculiar stars. Our detailed abundance analysis could provide information about the turbulence occurring in the outer stellar regions in Am stars with well determined age. In particular, our analysis can give constraints to define the depth of the zone mixed by turbulence, since it is the only parameter characterising turbulence \citep{richer2000}. A systematic abundance analysis of normal and peculiar stars in clusters could provide information on the origin of the mixing process and show if only turbulence is needed to explain abundance anomalies, or if other hydrodynamical processes occur. We tackle this problem using new and more precise new-generation spectrographs providing a wider wavelength coverage together with newer analysis codes and procedures (e.g. Least-Squares Deconvolution and synthetic line profile fitting instead of equivalent width measurements.) The observed stars, the instruments employed and the target selection are described in Sect.~\ref{observations}. The data reduction and a discussion of the continuum normalisation are provided in Sect.~\ref{reduction,norm}. In Sect.~\ref{LLmodels} and \ref{abundanceanalysis} we describe the models and the procedure used to perform the abundance analysis. Our results are summarised in Sect.~\ref{results}. Discussion and conclusions are given in Sect.~\ref{discussion} and \ref{conclusion} respectively. | 7 | 10 | 0710.0579 |
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0710 | 0710.2783_arXiv.txt | For the 2dFGRS we study the properties of voids and of fainter galaxies within voids that are defined by brighter galaxies. Our results are compared with simulated galaxy catalogues from the Millenium simulation coupled with a semianalytical galaxy formation recipe. We derive the void size distribution and discuss its dependence on the faint magnitude limit of the galaxies defining the voids. While voids among faint galaxies are typically smaller than those among bright galaxies, the ratio of the void sizes to the mean galaxy separation reaches larger values. This is well reproduced in the mock galaxy samples studied. We provide analytic fitting functions for the void size distribution. Furthermore, we study the galaxy population inside voids defined by objects with $B_J -5\log{h}< -20$ and diameter larger than 10 \hMpc. We find a clear bimodality of the void galaxies similar to the average comparison sample. We confirm the enhanced abundance of galaxies in the blue cloud and a depression of the number of red sequence galaxies. There is an indication of a slight blue shift of the blue cloud. Furthermore, we find that galaxies in void centers have higher specific star formation rates as measured by the $\eta$ parameter. We determine the radial distribution of the ratio of early and late type galaxies through the voids. We find and discuss some differences between observations and the Millenium catalogues. | The large-scale galaxy distribution is highly inhomogeneous. We observe groups, clusters and superclusters of galaxies and large voids. During last decades, much attention was paid on the analysis of bound structures as groups and clusters. Recently, new superclusters catalogues were constructed from the 2dFGRS and compared with large cosmological simulations \citep{Einasto07a, Einasto07b}. In a complement, there are large regions in the universe without bright galaxies, so called cosmic voids. Early on very large voids over 50 \hMpc diameter were found by \citet{Gregory78} and \citet{Kirshner81}. More common are voids with diameters of about 10 \hMpc that fill most of cosmic space. The explanation of such structures is not obvious. According to the standard paradigm of cosmological structure formation, negative potential wells from primordial inhomogeneities attract all matter in bound structures. In the same way, positive potential perturbations expel matter, but observed voids are too large for completely emptying. Therefore, in addition to the dilution of matter, the galaxy formation probability should be suppressed in underdense regions, cp. e.g. \citet{Lee98, Madsen98}. Recently, \citet{Sheth04} and \citet{Furlanetto06} applied these ideas within the excursion set formalism of gravitational instability. These analytical theories derived void size distributions that are peaked typically at diameters below 10 \hMpc which seem to be smaller than observed void sizes, cp. \citet{Mueller00}, and void sizes in CDM-simulations coupled with semianalytical galaxy formation models, \citet{Benson03}. Voids were routinely identified in all wide-field redshift surveys as the CfA \citep{deLapparent86, Vogeley94}, the SSRS2 \citep{Elad97}, the LCRS \citep{Mueller00, Arbabi02}, the IRAS-survey \citep{Elad00}, the 2dFGRS \citep{Hoyle04, Croton04, Patiri06a}, the SDSS \citep{Rojas04, Rojas05, Patiri06b}, and the DEEP2 survey with an analysis of voids up to redshift $z \approx 1$ \citep{Conroy05}. However, many void searches are only devoted to the identification of large voids, other void finders depend on special procedures as firstly identifying wall galaxies by an overdensity criterion and then looking for voids bounded by wall galaxies \citep{Elad97, Hoyle04}. Furthermore, the void search depends on the galaxy sample used for defining voids, in particular on the limiting magnitude of the galaxy sample. In an influential paper, \citet{Peebles01} derived from nearest neighbor statistics that galaxies of different brightness respect the same voids. He claimed that this contradicts the standard CDM scenario of galaxy and structure formation that seem to predict a hierarchy of galactic structures with smaller structure for fainter objects sitting in less massive dark matter halos, i.e also smaller voids for fainter objects. In a follow up theoretical study, \citet{Mathis02} showed from high-resolution simulation that voids defined by bright galaxies are also underdense in faint galaxies, i.e. that bright and faint galaxies respect similar voids. We want to take up this question since our earlier studies of voids in LCRS and in LCDM-simulations \citep{Mueller00, Arbabi02} showed a dependence of the void size distribution on the brightness limit of the galaxies under study and a characteristic void size scaling relation. \citet{Benson03} confirmed this scaling relation in simulated galaxy distributions, but the quantitative parameters were different. They suspected differences in the void search algorithms as reason, but we suspect that the effective 2-dimensional nature of the LCRS is the most likely cause. But their demand for using the same void search algorithm both in data and simulations seems a prerequisite for trustworth results. More recently, Colberg et al. (2007 in preparation) compared different void search algorithms and found that most proposed algorithm find comparable locations and sizes of large voids. This is very likely not the case for the large number of small voids that fill a significant part of space. Therefore we shall present in this study an comparable analysis of voids both in simulations and in the data that explores the detailed void size distribution in dependence on the faint brightness limit of the galaxies defining the voids. We shall use for our study the 2dFGRS \citep{Cole05} thereby coming back to the property of the self-similarity of the void statistics. It tells that the void size distribution depends on the mean galaxy separation, in such a way that brighter galaxies define larger voids than fainter ones. Even if voids in the 2dFGRS were previously analyzed \citep{Hoyle04, Patiri06a}, this concerned mainly large voids and not the detailed void statistics proposed by us previously. \citet{Croton04} provided a detailed study of the void probability distribution for the 2dFGRS which is related to the void size distribution but it provides an different statistics. Essentially the void probability distribution is a weighted sum over the void size distribution \citep{Otto86}. The 2dFGRS is a densely sampled survey with a compact survey geometry. This is of advantage for the question of the dependence of the void sizes on the galaxy magnitudes defining voids. We shall derive phenomenological fits to the void size distribution that will be compared with simulation results. Furthermore, we shall take up the question of the faint galaxies within voids. Thereby we cut both questions, the matter content inside large underdense regions in the universe, and the change of galaxy properties. The color distribution of galaxies in the 2dFGRS employs SuperCOSMOS data \citep{Hambly01} for the $R$-band. We will find a clear bimodality in the void galaxies so far only studied in detail for the SDSS \citep{Rojas05, Patiri06b} but not yet for the 2dFGRS. We shall evaluate our void analysis with model galaxy samples constructed from the Millenium simulation of \citet{Springel05} and from semianalytical galaxy formation theory applied to the numerical merger trees \citep{Croton06}. We analyzed specific galaxy properties within voids and found results that can be qualitatively described by the model samples. A quantitative comparison of the galaxy color distribution and the star formation efficiency hints at certain differences between observations and simulations. Tentatively we connect it with specific environmental properties of galaxy formation in underdense regions. In particular, major mergers, galaxy harassment, tidal and ram-pressure stripping will not be as effective there as in more dense regions of the universe \citep{Avila05, Maulbetsch06}. The outline of the paper is as follows: First we describe the galaxy extraction from the 2dFGRS, and in Section 3 we provide some details of the galaxy mock data. In Section 4 we describe our void search algorithm and in Section 5 we provide our results. Section 6 is devoted to a discussion and in the final Section we draw our conclusions. | In this paper we have studied properties of voids in the 2dFGRS and compared them with those predicted by semi-analytical models of galaxy formation. In particular, we were interested in the distribution of void sizes in the galaxy distribution as a function of the faint limiting magnitude of the sample. We established an almost linear dependence of the void size distribution on the mean separation $\lambda$ of the galaxies in the sample. A similar relation was found for the magnitude dependence of the correlation length of galaxies and galaxy groups \citep{Bahcall92, Yang05}. It is a natural consequence of the halo model of gravitational clustering and the resulting void statistics \citep{Tinker06}. In addition to the self-similarity, we have found that the voids among faint galaxies extend to relatively larger scales when devided by the mean void size. In the scaled radius $R/\lambda$, the voids size distribution of faint galaxies has a longer tail at large radii than those for brighter galaxies. This indicates that the faint galaxies trace the general spatial distribution of the cosmic web of brighter galaxies. For the 2dFGRS, we find more large voids at larger distances from the observer. This is due to the conelike structure of the survey. In our mock sample, we use exactly the same geometry of the observed samples. Thereby we found that the number density of the largest voids can be underestimated by up to 20\%. We confirm the previous findings by \citet{Grogin00, Hogg04, Rojas05, Croton05, Patiri06b}, that in lower-density regions, the galaxy population is dominated by blue actively star forming galaxies. We fitted the $B_J-R$ color distribution of 2dFGRS void galaxies by double Gaussian distributions. There are significantly more void galaxies in the blue cloud than in the general field, and the red sequence is strongly suppressed. In addition, we have found the indication that the blue population of galaxies in void centers is slightly {\it bluer} than the field population, not just more numerous. The shift is only minor and almost of the same order as the maximal quoted uncertainty by \cite{Cole05}, so it remains unclear how robust is this effect. It is not reproduced by the semi-analytical models for galaxy formation of the Millenium simulation. Also the red sequence in the Millenium catalogue has a significantly tighter spread. This may due to a too sharp cutoff of the star formation activity in the semianalytical models which seems to be unrealistic. The radial distribution of the ratio of early and late type galaxies inside voids coincides quantitatively for the 2dFGRS and the Millenium catalogue, but in the data, the fraction of E-galaxies stays smaller (and the fraction of S-galaxies larger) further outwards beyond the boundary of the voids. The abundance of star-forming galaxies is higher in the 2dFGRS voids than in the semianalytical model. | 7 | 10 | 0710.2783 |
0710 | 0710.3785_arXiv.txt | % In order to distinguish between the various components of massive star forming regions (i.e. infalling, outflowing and rotating gas structures) within our own Galaxy, we require high angular resolution observations which are sensitive to structures on all size scales. To this end, we present observations of the molecular and ionized gas towards massive star forming regions at 230 GHz from the SMA (with zero spacing from the JCMT) and at 22 and 23 GHz from the VLA at arcsecond or better resolution. These observations (of sources such as NGC7538, W51e2 and K3-50A) form an integral part of a multi-resolution study of the molecular and ionized gas dynamics of massive star forming regions (i.e. Klaassen \& Wilson 2007). Through comparison of these observations with 3D radiative transfer models, we hope to be able to distinguish between various modes of massive star formation, such as ionized or halted accretion (i.e Keto 2003 or Klaassen et al. 2006 respectively). | At the large distances to massive star forming regions, we do not yet quite have the resolution necessary to determine whether or not they can form via a scaled up version of the processes responsible for the formation of low mass stars. The energetics (i.e. luminosity, outflow energy and accretion rates) are orders of magnitude higher than those seen in low mass star forming regions (e.g. Arce et al. 2007), and the outward radiation and thermal pressures produced when the massive star begins ionizing its surroundings are enormous. Yet, collimated outflows and both ionized and molecular infall signatures are present in high mass star forming regions just as they are in their lower mass counterparts. If we assume that massive stars can form in a single accretion event similar to their low mass counterparts, how then can accretion continue beyond the formation of the H{\sc II} region? Does accretion continue in an ionized form? Does it stop early in the star's evolution? Does it continue through a disk (not yet observed around forming O stars)? G10.62-0.38 (hereafter G10.6) is a well studied ultracompact H{\sc II} (UCH{\sc II}) region (Ho et al. 1981, Keto et al. 1988, etc) through which the surrounding gas is falling in, becoming ionized, and continuing to fall in towards the central star. The high resolution observations of Keto \& Wood (2006) in H66$\alpha$ taken at the VLA show an infall signature in their PV diagram, and so it appears as though accretion onto the protostar is, in this case, ionized and is continuing beyond the formation of the UCH{\sc II} region. G5.89-0.39 (hereafter G5.89) is the site of an UCH{\sc II} region (the archetypal shell UHCII region from Wood \& Churchwell 1989) and multiple outflows depending on the molecular tracer observed. This region appears to be the formation site for a number of massive stars including the O5V star which is believed to power the H{\sc II} region (Feldt et al. 2003) and the 1.3 mm continuum source believed to be powering the SiO outflow detected by Sollins et al. (2004). Recent maser emission studies of this region (Stark et al. 2007, Fish et al. 2005) also suggest that the two sources are independent. The star identified in the mid-infrared by Feldt et al. (2003), corresponding to `G5.89 center' in Stark et al (2007), appears to be less embedded than the Sollins et al. (2004) 1.3 mm source, suggesting it is possibly more evolved than the 1.3 mm source which drives the accelerating SiO outflow. To date, there are no published CO maps of this region at high enough resolution to determine which of these two sources (separated by $\sim2''$) lies at the center of the large scale decelerating outflow. However, the velocity gradient in the 1667 MHz maser emission towards G5.89 center (Stark et al. 2007) is consistent with the direction of the large scale east-west CO outflow observed by Klaassen et al. (2006) and Watson et al. (2007). Klaassen et al. (2006) suggested that accretion onto the source powering the large scale outflow may have halted at the onset of the H{\sc II} region. This conclusion appears consistent with the O5 star being responsible for the H{\sc II} region, `G5.89 center' maser emission and the large scale CO outflow, while the 1.3 mm continuum source from Sollins et al. (2004) is a younger protostar responsible for the SiO outflow. These studies of G10.6 and G5.89 suggest that high neutral accretion rates (up to 10$^{-2}$ M$_{\odot}$/yr, Edgar \& Clarke 2003) and continued ionized accretion can both account for the formation of massive stars. Yet, these conclusions are based on the in depth studies of two individual regions, and from these two regions, we cannot make conclusive statements about the general nature of massive star formation. In order to determine whether there is a preferred method, we need to observe these processes in a number of regions, and build up statistics about whether/how accretion onto a massive protostar occurs beyond the onset of dynamical expansion in the UCH{\sc II} regions. | 7 | 10 | 0710.3785 |
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0710 | 0710.3099_arXiv.txt | Short--hard Gamma Ray Bursts (SGRBs) are currently thought to arise from gravitational wave driven coalescences of double neutron star systems forming either in the field or dynamically in globular clusters. For both channels we fit the peak flux distribution of BATSE SGRBs to derive the local burst formation rate and luminosity function. We then compare the resulting redshift distribution with {\it Swift} 2-year data, showing that both formation channels are needed in order to reproduce the observations. Double neutron stars forming in globular clusters are found to dominate the distribution at $z\lsim 0.3$, whereas the field population from primordial binaries can account for the high--$z$ SGRBs. This result is not in contradiction with the observed host galaxy type of SGRBs. | The afterglows of several Short Gamma Ray Bursts (SGRBs) have recently been localised on the sky with {\it Swift} (Gehrels et al. 2004) and {\it HETE-II} (Lamb et al. 2004), allowing for the determination of their redshift and host galaxy (Nakar 2007; Berger et al. 2007a, and references therein). For SGRBs, the data give mean redshift $z\sim 0.6,$ and early as well as late type galaxy hosts. The current idea on the nature of the SGRBs is that they arise from gravitational wave driven mergers of neutron star-neutron star and/or neutron star-black hole binaries (Belczynski et al. 2006; Nakar 2007). Double neutron star binaries (DNSs) are observed in the field of the Milky Way and in globular clusters, so they can form through two different channels: (i) evolution of primordial massive star binaries, formed in the galactic field (field scenario; Narayan, Paczynski \& Piran 1992), and (ii) dynamical formation\footnote{Formation of primordial DNSs may not be efficient in globular clusters (Ivanova et al 2008). }, through three-body interactions in globular clusters (GCs) involving the exchange of a light star companion of a neutron star with an incoming isolated neutron star (GC scenario; Grindlay, Portegies Zwart \& Mc Millan 2006). In this Letter, we derive the local SGRB formation rate and luminosity function by fitting the peak flux distribution of the large sample of BATSE SGRBs (Paciesas et al. 1999) exploring both formation scenarios. We compare the model results with the redshift distribution of the sample of {\it Swift} SGRBs in the first two years of operation, with the aim of disentangling the relative importance of the two SGRB populations at different cosmic epochs. | We have shown that current {\it Swift} data seem to point towards a dual nature of SGRB formation: low--$z$ SGRBs may arise from the coalescence of DNSs forming in GCs through dynamical collisions, whereas high--$z$ bursts can represent the end--result of DNS mergers born in the field from primordial massive binary stars. A way to falsify this bimodality is by considering the correlation between SGRBs and the host galaxy type (Gam-Yam et al. 2005; Zheng \& Ramirez-Ruiz 2007; Berger et al. 2007b; Shin \& Berger 2007). We expect that SGRBs from the GC channel should be found preferentially in early type galaxies where the bulk of GCs resides, although we can not completely exclude a few in spirals. On the contrary, field SGRBs can be hosted in both early and late type galaxies given the wide distribution of delay times (Section 2). Since the dynamical channel produces SGRBs at low--$z$, we expect to find a excess of SGRBs in early type galaxies below $z<0.3$. At present it is still premature to exploit this issue in a statistically meaningful manner and we may just confine ourselves to the analysis of SGRBs in the current limited sample. Up to now, four SGRBs have been localised in early-type galaxies (GRB~050509B, 050724, 050813, and 060502B), whereas six have secure late-type hosts (GRB~051221A, 060801, 061006, 061210, and 061217, Berger 2007; {\it HETE-II} GRB~050709 at $z=0.161$; Fox et al. 2005; Covino et al. 2006). Fig.~2 shows the fraction of SGRB hosts in early (shaded area) and late (white area) below $z=0.3$ and above, respectively. We note indeed that low--$z$ SGRBs are associated preferentially to early-type galaxies while the high--$z$ ones are located in late-type. The large off--set of GRB~050509B and 060502B may support further the association of these SGRBs with DNSs formed dynamically in GCs. \begin{figure} \begin{center} \centerline{\psfig{figure=host.ps,height=8cm}} \caption{Fraction of SGRBs hosted in early--type (shaded area) and late--type (white area) galaxies below $z=0.3$ and above.} \end{center} \end{figure} | 7 | 10 | 0710.3099 |
0710 | 0710.1113_arXiv.txt | We investigate torsional Alfv\'{e}n oscillations of relativistic stars with a global dipole magnetic field, via two-dimensional numerical simulations. We find that a) there exist two families of quasi-periodic oscillations (QPOs) with harmonics at integer multiples of the fundamental frequency, b) the lower-frequency QPO is related to the region of closed field lines, near the equator, while the higher-frequency QPO is generated near the magnetic axis, c) the QPOs are long-lived, d) for the chosen form of dipolar magnetic field, the frequency ratio of the lower to upper fundamental QPOs is $\sim 0.6$, independent of the equilibrium model or of the strength of the magnetic field, and e) within a representative sample of equations of state and of various magnetar masses, the Alfv\'{e}n QPO frequencies are given by accurate empirical relations that depend only on the compactness of the star and on the magnetic field strength. The lower and upper QPOs can be interpreted as corresponding to the edges or turning points of an Alfv\'{e}n continuum, according to the model proposed by Levin (2007). Several of the low-frequency QPOs observed in the X-ray tail of SGR 1806-20 can readily be identified with the Alfv\'{e}n QPOs we compute. In particular, one could identify the 18Hz and 30Hz observed frequencies with the fundamental lower and upper QPOs, correspondingly, while the observed frequencies of 92Hz and 150Hz are then integer multiples of the fundamental upper QPO frequency (three times and five times, correspondingly). With this identification, we obtain an upper limit on the strength of magnetic field of SGR 1806-20 (if is dominated by a dipolar component) between $\sim3$ and $7\times 10^{15}$G. Furthermore, we show that an identification of the observed frequency of 26Hz with the frequency of the fundamental torsional $\ell=2$ oscillation of the magnetar's crust is compatible with a magnetar mass of about $1.4$ to 1.6$M_\odot$ and an EOS that is very stiff (if the magnetic field strength is near its upper limit) or moderately stiff (for lower values of the magnetic field). | \label{sec:Intro} The phenomenon of Soft Gamma Repeaters (SGRs) may allow us in the near future to determine fundamental properties of strongly magnetized, compact stars. Already, there exist at least two sources in which quasi-periodic oscillations (QPOs) have been observed in their X-ray tail, following the initial discovery by \cite{Israel2005}, see \cite{WS2006} for a recent review. The frequency of many of these oscillations is similar to what one would expect for torsional modes of the solid crust of a compact star. This observation is in support of the proposal that SGRs are magnetars (compact objects with very strong magnetic fields) \citep{DT1992}. During an SGR event, torsional oscillations in the solid crust of the star could be excited \citep{Duncan1998}, leading to the observed frequencies in the X-ray tail. However, not all of the observed frequencies fit the above picture. For example, the three lowest observed frequencies for SGR 1806-20 are 18, 26, 30Hz. Only one of these could be the fundamental, $\ell=2,m=0$ torsional frequency of the crust, as the first overtone has a much higher frequency. \cite{Levin2006} stressed the importance of crust-core coupling by a global magnetic field and of the existence of an Alfv\'en continuum, while \citet{GSA2006} considered model with simplified geometry, in which Alfv\'{e}n oscillations form a discrete spectrum of normal modes, that could be associated with the observed low-frequency QPOS. In \cite{Levin2007}, the existence of a continuum was stressed further and it was shown that the edges or turning points of the continuum can yield long-lived QPOs. In addition, numerical simulations showed that drifting QPOs within the continuum become amplified near the frequencies of the crustal normal modes. Within this model, Levin suggested a likely identification of the 18Hz QPO in SGR 1806-20 with the lowest frequency of the MHD continuum or its first overtone. The above results were obtained in toy models with simplified geometry and Newtonian gravity. In this Letter, we perform two-dimensional numerical simulations of linearized Alfv\'{e}n oscillations in magnetars. Our model improves on the previously considered toy models in various ways: relativistic gravity is assumed, various realistic equations of state (EOS) are considered and a consistent dipolar magnetic field is constructed. We do not consider the presence of a solid crust, but only examine the response of the ideal magnetofluid to a chosen initial perturbation. Spherical stars have generally two type of oscillations, {\it spheroidal} with polar parity and {\it toroidal} with axial parity. The observed QPOs in SGR X-ray tails may originate from toroidal oscillations, since these could be excited more easily than poloidal oscillations, because they do not involve density variations. In Newtonian theory, there have been several investigations of torsional oscillations in the crust region of neutron stars (see e.g., \cite{Lee2007} for reference). On the other hand, only few studies have taken general relativity into account \citep{Messios2001,Sotani2007a,Sotani2007b,SA2007,Vavoulidis2007}. SGRs produce giant flares with peak luminosities of $10^{44}$ -- $10^{46}$ erg/s, which display a decaying tail for several hundred seconds. Up to now, three giant flares have been detected, SGR 0526-66 in 1979, SGR 1900+14 in 1998, and SGR 1806-20 in 2004. The timing analysis of the latter two events revealed several QPOs in the decaying tail, whose frequencies are approximately 18, 26, 30, 92, 150, 625, and 1840 Hz for SGR 1806-20, and 28, 53, 84, and 155 Hz for SGR 1900+14, see \cite{WS2006}. In \cite{Sotani2007a} (hereafter Paper~I), it was suggested that some of the observational data of SGRs could agree with the crustal torsional oscillations, if, e.g., frequencies lower than 155 Hz are identified with the fundamental oscillations of different harmonic index $\ell$, while higher frequencies are identified with overtones. However, in Paper~I and above, it will be quite challenging to identify all observed QPO frequencies with only crustal torsional oscillations. For example, it is difficult to explain all of the frequencies of 18, 26 and 30 Hz for SGR 1806-20 with crustal models, because the actual spacing of torsional oscillations of the crust is larger than the difference between these two frequencies. Similarly, the spacing between the 625Hz and a possible 720Hz QPO in SGR 1806-20 may be too small to be explained by consecutive overtones of crustal torsional oscillations. One can notice, however, that the frequencies of 30, 92 and 150 Hz in SGR 1806-20 are in near {\it integer ratios}. As we will show below, the numerical results presented in this Letter are compatible with this observation, as we find two families of QPOs (corresponding to the edges or turning points of a continuum) with harmonics at near integer multiples. Furthermore, our results are compatible with the ratio of 0.6 between the 18 and 30Hz frequencies, if these are identified, as we suggest, with the edges (or turning points) of the Alfv\'{e}n continuum. With this identification, we can set an upper limit to the dipole magnetic field of $\sim3$ to $~7\times 10^{15}$G. If the drifting QPOs of the continuum are amplified at the fundamental frequency of the crust, and the latter is assumed to be the observed 26Hz for SGR 1806-20, then our results are compatible with a magnetar mass of about $1.4$ to 1.6$M_\odot$ and an EOS that is very stiff (if the magnetic field strength is near its upper limit) or moderately stiff (for lower values of the magnetic field). Unless otherwise noted, we adopt units of $c=G=1$, where $c$ and $G$ denote the speed of light and the gravitational constant, respectively, while the metric signature is $(-,+,+,+)$. \section[]{Numerical Setup} \label{sec:II} The general-relativistic equilibrium stellar model is assumed to be spherically symmetric and static, i.e. a solution of the well-known TOV equations for and perfect fluid and metric described the line element \begin{equation} ds^2 = -e^{2\Phi(r)}dt^2 + e^{2\Lambda(r)}dr^2 + r^2(d\theta^2 + \sin^2\theta d\phi^2). \end{equation} We neglect the influence of the magnetic field on the structure of the star, since the magnetic field energy, ${\cal E_M}$, is orders of magnitudes smaller than the gravitational binding energy, ${\cal E_G}$, for magnetic field strengths considered realistic for magnetars, ${\cal E_M}/{\cal E_G}\approx 10^{-4}(B/(10^{16} {\rm G}))^2$. For simplicity, we assume that the magnetic field is a pure dipole (toroidal magnetic fields will be treated elsewhere, see \cite{Sotani2007c}. Details on the numerical method for constructing the magnetic field, as well as representative figure of the magnetic field lines, can be found in Paper~I. MHD oscillations of the above equilibrium model are described by the linearized equations of motion and the magnetic induction equations, presented in detail in Papers~I and II (we neglect perturbations in the spacetime metric, as these couple weakly to toroidal modes in a spherically symmetric background). The perturbative equations presented in Papers~I and II can readily be converted from an eigenvalue problem to the form of a two-dimensional time-evolution problem, by defining a displacement ${\cal Y}(t,r,\theta)$ due to the toroidal motion (the coefficient of shear viscosity is set to $\mu=0$, as we neglect the presence of a solid crust in the present work). The contravariant {\it coordinate component} of the perturbed four-velocity, $\delta u^{\phi}$, is then related to time derivative of ${\cal Y}$ through \begin{equation} \delta u^{\phi} = e^{-\Phi}\partial_t {\cal Y}(t,r,\theta). \end{equation} The two-dimensional evolution equation for ${\cal Y}(t,r,\theta)$ is \begin{eqnarray} {\cal A}_{tt}\frac{\partial^2 {\cal Y}}{\partial t^2} &=& {\cal A}_{20}\frac{\partial^2 {\cal Y}}{\partial r^2} + {\cal A}_{11} \frac{\partial^2 {\cal Y}}{\partial r \partial \theta} + {\cal A}_{02} \frac{\partial^2 {\cal Y}}{\partial \theta^2} \nonumber \\ &+& {\cal A}_{10} \frac{\partial {\cal Y}}{\partial r} + {\cal A}_{01} \frac{\partial {\cal Y}}{\partial \theta} + \varepsilon_D {\cal D}_4 {\cal Y}, \label{eq2D} \end{eqnarray} where ${\cal A}_{tt}$, ${\cal A}_{20}$, ${\cal A}_{11}$, ${\cal A}_{02}$, ${\cal A}_{10}$, and ${\cal A}_{01}$ are functions of $r$ and $\theta$, given by \begin{eqnarray} {\cal A}_{tt} &=& \left[\epsilon + p + \frac{{a_1}^2}{\pi r^4} \cos^2\theta + \frac{{{a_1}'}^2}{4\pi r^2} e^{-2\Lambda}\sin^2\theta \right] e^{-2(\Phi - \Lambda)}, \\ {\cal A}_{20} &=& \frac{{a_1}^2}{\pi r^4}\cos^2\theta , \\ {\cal A}_{11} &=& -\frac{a_1 {a_1}'}{\pi r^4}\sin\theta \cos\theta, \\ {\cal A}_{02} &=& \frac{{{a_1}'}^2}{4\pi r^4}\sin^2\theta, \\ {\cal A}_{10} &=& \left(\Phi' - \Lambda' \right) \frac{{a_1}^2}{\pi r^4} \cos^2\theta + \frac{a_1 {a_1}'}{2\pi r^4} \sin^2\theta, \\ {\cal A}_{01} &=& \left[\frac{a_1}{\pi r^4}\left(2\pi j_1 - \frac{a_1}{r^2}\right)e^{2\Lambda} + \frac{3{{a_1}'}^2}{4\pi r^4} \right]\sin\theta\cos\theta, \end{eqnarray} and where $a_1(r)$ and $j_1(r)$ are the radial components of the electromagnetic four-potential and the four-current, respectively. In the above equations, a prime denotes a partial derivative with respect to the radial coordinate. In our numerical scheme, we employ the 2nd-order, iterative Crank-Nicholson scheme. A numerical instability, which sets in after many oscillations, was treated by adding a 4th-order Kreiss-Oliger dissipation term \citep{KO1973}, shown as $\varepsilon_D {\cal D}_4 {\cal Y}$ in Equation (\ref{eq2D}) above. We experimented with various values of the dissipation coefficient $\varepsilon_D$ and found the evolution to be stable for values as small as a few times $10^{-5}$. We verified that in this limit, the solution becomes independent of the strength of the numerical dissipation, as the ${\cal D}_4$ Kreiss-Oliger dissipation operator introduces an error of higher order than the 2nd-order iterative Crank-Nicholson scheme. The numerical grid we use is equidistant, covering only the interior of the star, with (typically) 50 radial zones and 40 angular zones (we also compared our results to simulations with 100x80 points). The boundary conditions are: (a) ${\cal Y}=0$ at $r=0$ (regularity), (b) ${\cal Y}_{,r}=0$ at $r=R$ (vanishing traction) (c) ${\cal Y}_{,\theta}=0$ at $\theta=0$ (axisymmetry) and (d) ${\cal Y}=0$ at $\theta=\pi/2$ (equatorial plane symmetry of the $\ell=2$ initial data). We obtain the same results if you use a grid that extends to $\theta=\pi$. We have also evolved initial data with $\ell=3$, for which the appropriate boundary condition at $\theta=\pi/2$ is ${\cal Y}_{,\theta}=0$. \section[]{Alfv\'{e}n QPOs} \label{sec:III} \begin{figure} \begin{center} \includegraphics[width=55mm]{Fig-fft-e0001-dy1.eps} \end{center} \caption{FFT of the MHD oscillations at $\theta\sim 0$ (magnetic axis). Three lines in each figure correspond to different radial positions $r\sim 0$, $R/2$, and $R$). A fundamental ($n=0$) QPO and several overtones (nearly integer multiples) are clearly present.} \label{Fig:FFT} \end{figure} \begin{figure} \begin{center} \includegraphics[width=55mm]{Fig-fft-e0001-dy3.eps} \end{center} \caption{Same as Fig. \ref{Fig:FFT}, but at $\theta\sim \pi/2$ (magnetic equator). A second family of QPOs is present. Arrows indicate several continuous parts, of which only the first is distinct from the others, which partially overlap.} \label{Fig:FFT2} \end{figure} \begin{figure} \begin{center} \begin{tabular}{cc} \includegraphics[width=36mm]{e000002-dy-u0-unitbase.eps} & \includegraphics[width=36mm]{e000002-dy-u1-unitbase.eps} \\ \includegraphics[width=36mm]{e000002-dy-l0-unitbase.eps} & \includegraphics[width=36mm]{e000002-dy-l1-unitbase.eps} \\ \end{tabular} \end{center} \caption{Distribution of effective amplitude of several Alfv\'{e}n QPOs (see text for details). The grayscale map varies from white (zero amplitude) to black (maximum amplitude).} \label{Fig:Effective} \end{figure} As a fist equilibrium model, we examined a polytropic stellar model with $\Gamma=2$, whose mass and radius are $M=1.4M_{\odot}$ and $R=14.16$ km, respectively, with a dipole magnetic field strength of $B\equiv B_\mu = 4\times 10^{15}$ G, where $B_{\mu}$ is a typical strength. As initial data, we use the numerical eigenfunction of the $\ell=2,m=0$ fundamental mode, for the truncated system presented in Paper~II. The grid size is $100 \times 80$ and $\varepsilon_D=10^{-3}$, while the total simulation time is 2s. By computing the FFT of $\partial_t{\cal Y}$ (which is proportional to $\delta u^\phi$) at various points inside the star, we obtain the following results: Examining the FFT at three different radial locations ($r\sim 0, R/2, R$) at $\theta \sim 0$ (magnetic axis), Fig. \ref{Fig:FFT}, we observe a number of narrow frequency peaks. The strength of the peaks is several orders of magnitude larger than the FFT continuum. The overtones ($n>0$) are nearly {\it integer multiples} of the fundamental frequency ($n=0)$. The corresponding FFTs at the same radial points, but at $\theta \sim \pi/2$ (magnetic equator) are shown in Fig. \ref{Fig:FFT2}, where, in addition to the family of frequency peaks observed in Fig. \ref{Fig:FFT}, one can also observe a second family of frequency peaks, mainly at $r\sim R$, for which the fundamental frequency is at a ratio of 0.6 with respect to the fundamental frequency in Fig. \ref{Fig:FFT}. The fundamental frequency of this family, as well as the first overtone (again an integer multiple) are clearly visible, while the other overtones seem to be buried inside a continuum of other frequencies. Using Levin's toy model for the Alfv\'{e}n continuum in magnetars, the above numerical results can be interpreted as follows: The fundamental frequency peaks of the two families are QPOs generated by the edges or turning points of an Alfv\'{e}n continuum (henceforth we call these two frequencies as the fundamental lower and upper QPOs and denote them as $L_0$ and $U_0$). We denote the overtones (which are nearly integer multiples) as $L_n$ and $U_n$, where $n>0$. For the chosen form of the magnetic field, the extent of the first continuum, between the fundamental $L_0$ and $U_0$ QPOs is at lower frequencies than the continuous frequency intervals corresponding to the overtones, which all overlap partially. Because of the Alfv\'{e}n continuum, the phase of the oscillations at the QPO frequencies is not constant, but varies throughout the star. Since the magnetic field is axisymmetric, the axis is both a turning point and edge of the continuum. Near the magnetic axis, the background magnetic field varies slowly, which allows for the upper QPOs to remain nearly coherent for a large number of oscillations. Since $\delta u^\phi$ is only a coordinate component, we also obtained FFTs of $r \sin\theta \partial_t {\cal Y}$ (which is proportional to the physical velocity in a unit basis) at all grid points. In two-dimensional simulations of fluid modes in non-magnetized stars, it has been shown that the amplitude of the FFT of physical variables at every point in the numerical grid, at a chosen normal mode frequency, is correlated to the shape of the eigenfunction of a given mode \citep{Stergioulas2004}. Similarly, we use the magnitude of the FFT of $r \sin\theta \partial_t {\cal Y} $ at every point in the grid to define an {\it effective amplitude} for each QPO (which, of course, will be time-varying, due to the absorption of individual modes in the continuum). In Fig. \ref{Fig:Effective}, we show the effective amplitude of the fundamental upper QPO, $U_0$ and its first overtone $U_1$ as well as for the corresponding lower QPOs $L_0$ and $L_1$, obtained after a simulation with a $50\times 40$ grid, $\varepsilon_D=2\times 10^{-5}$ and a duration of 2.1s. The effective amplitude for $U_0$ has a maximum near the magnetic axis, while there are also {\it nodal lines} along certain magnetic field lines. For the overtone, $U_1$, there exists an additional nodal line, starting perpendicular near the {\it magnetic axis}, which divides the region of maximum amplitude near the magnetic axis roughly in half. Similarly (not shown here), each successive overtone corresponds to an additional ``horizontal'' nodal line, dividing the region of maximum amplitude near the magnetic axis into roughly equidistant parts. This agrees well with the fact that the frequencies of the overtones are nearly integer multiples of the fundamental frequency. The effective amplitude of the fundamental lower QPO, $L_0$, is practically limited to a region that is only somewhat larger than the region of {\it closed magnetic field lines} inside the star. For the first overtone $L_1$, a nodal line divides the region of maximum amplitude into two parts. In Fig. \ref{Fig:evol1} we show the evolution of $\partial_t {\cal Y}$ at the location inside the star where the effective amplitude of the fundamental upper QPO, $U_0$ attains its maximum value. It is evident that the QPO is long-lived, since the amplitude of the oscillations barely diminishes with time. \begin{figure} \begin{center} \includegraphics[width=70mm]{50-40-e000002-dy-m6-x34.eps} \end{center} \caption{Time evolution of $\partial_t {\cal Y}$ at the location inside the star where the effective amplitude of the fundamental upper QPO, $U_0$ attains its maximum value. The amplitude of the oscillations barely diminish with time} \label{Fig:evol1} \end{figure} \begin{table*} \centering \caption{Frequencies of lower and upper Alfv\'{e}n QPOs and their ratios, for a representative sample of equilibrium models, constructed with various EOSs and masses and for a magnetic field strength of $B=B_\mu$ (see text for details).} \label{Tab:eos1} \begin{tabular}{@{}lcccccccc@{}} \hline Model & M/R & $f_{L_0}$ (Hz) & $f_{U_0}$ (Hz) & ratio & $f_{L_1}$ (Hz) & $f_{U_1}$ (Hz) & ratio & $f_{U_2}$ (Hz) \\ \hline A+DH$_{14}$ & 0.218 & 15.4 & 25.0 & 0.616 & 30.7 & 49.4 & 0.621 & 74.4\\ A+DH$_{16}$ & 0.264 & 11.7 & 18.3 & 0.639 & 23.5 & 35.7 & 0.658 & 54.0\\ WFF3+DH$_{14}$ & 0.191 & 17.9 & 29.8 & 0.601 & 36.2 & 59.2 & 0.611 & 89.8\\ WFF3+DH$_{18}$ & 0.265 & 11.7 & 18.0 & 0.650 & 23.5 & 35.5 & 0.662 & 53.3\\ APR+DH$_{14}$ & 0.171 & 20.4 & 34.1 & 0.598 & 41.3 & 68.6 & 0.602 & 104.6\\ APR+DH$_{20}$ & 0.248 & 12.8 & 20.6 & 0.621 & 26.0 & 40.3 & 0.645 & 61.0\\ L+DH$_{14}$ & 0.141 & 23.7 & 40.8 & 0.581 & 47.5 & 81.6 & 0.582 & 123.8\\ L+DH$_{20}$ & 0.199 & 16.4 & 27.8 & 0.590 & 33.1 & 54.7 & 0.605 & 82.6\\ \hline \end{tabular} \end{table*} \begin{figure} \begin{center} \includegraphics[width=55mm]{f-MoR-fit.eps} \end{center} \caption{Quadratic fits in terms of the compactness of the star, $M/R$, of the lower and upper fundamental Alfv\'{e}n QPO frequencies, obtained for a representative sample of equilibrium models with various EOSs and masses. The magnetic field was set to $B=B_\mu$. } \label{Fig:fit} \end{figure} | \label{sec:V} We have already verified that our main QPO frequencies agree with frequencies obtained with an independent, fully nonlinear numerical code \citep{CerdaDuran2007}, for the same initial model. We caution, however, that we have not yet considered the crust-core interaction, different magnetic field topologies or the coupling to the exterior magnetosphere. These effects have to be taken into account and already \cite{Sotani2007c}, find that the observed QPOs could lead to constraints on the magnetic field topology. To complete the picture, a three-dimensional numerical simulation, that includes a proper coupling of the crust to the MHD interior and to the exterior magnetosphere will be required and our current results provide a good starting point. Extensive details of our computations will be presented in \cite{Sotani2007d}. | 7 | 10 | 0710.1113 |
0710 | 0710.3428_arXiv.txt | We explore a simple model for the high luminosity of SN~2006gy involving photon diffusion of shock-deposited thermal energy. The distinguishing property of the model is that the large ``stellar'' radius of $\sim$160 AU required to prevent adiabatic losses is not the true stellar radius, but rather, the radius of an opaque, unbound circumstellar envelope, created when $\sim$10 M$_{\odot}$ was ejected in the decade before the supernova in an eruption analogous to that of $\eta$~Carinae. The supernova light is produced primarily by diffusion of thermal energy following the passage of the blast wave through this shell. This model differs from traditional models of supernova debris interacting with external circumstellar matter (CSM) in that here the shell is optically thick and the escape of radiation is delayed. We show that any model attempting to account for SN~2006gy's huge luminosity with radiation emitted by ongoing CSM interaction fails for the following basic reason: the CSM density required to achieve the observed luminosity makes the same circumstellar envelope opaque ($\tau\ga$300), forcing a thermal diffusion solution. In our model, the weaker CSM interaction giving rise to SN~2006gy's characteristic Type IIn spectrum and soft X-rays is not linked to the power source of the visual continuum; instead, it arises {\it after} the blast wave breaks free of the opaque shell into the surrounding wind. While a simple diffusion model can explain the gross properties of the early light curve of SN~2006gy, it predicts that the light curve must plummet rapidly at late-times, unless an additional power source is present. | The recent extremely luminous supernova (SN) 2006gy has challenged our understanding of the deaths of massive stars and may represent a new class of explosions. It was discovered by Quimby (2006) and was subsequently studied in detail by Ofek et al.\ (2007) and Smith et al.\ (2007). It is a Type IIn event, with strong narrow lines of H in its spectrum, although the details of its light curve and spectrum are unlike most other members of this class (e.g., Filippenko 1997). Its enormous total radiated energy of $>$10$^{51}$ ergs, the late ($\sim$70 days) peak of the light curve, and the observed constant expansion speed of $v_{exp}\simeq$4000 km s$^{-1}$ (Smith et al.\ 2007) place important constraints on the model. SN~2006gy exploded in the host galaxy NGC~1260, which is a peculiar S0/Sa galaxy with ongoing star formation (see Smith et al.\ 2007 and Ofek et al.\ 2007). The optical displays of all Type Ia SNe and most Type II SNe are believed to be dominated by the radioactive decay sequence $^{56}$Ni$ \rightarrow ^{56}$Co$ \rightarrow ^{56}$Fe. If that is also the case for SN~2006gy, a mass M($^{56}$Ni)$\approx$10 M$_{\odot}$ is required to produce the observed luminosity (Smith et al.\ 2007). Not surprisingly, SN~2006gy has prompted varying suggestions to account for its extreme energy budget. Ofek et al.\ (2007) proposed a SN~Ia exploding in a dense H-rich circumstellar medium (CSM) where continual interaction with an external CSM was the power source for the luminosity. They suggested that the dense CSM was probably the result of binary star common envelope ejection, and that the extreme energy budget may need to draw upon a super-Chandrasekhar Type~Ia explosion or perhaps a massive star explosion. Smith et al.\ (2007) argued that many different spectral properties of SN~2006gy and its CSM were incompatible with a Type~Ia explosion, but were entirely consistent with the known properties of some very massive stars. To account for the luminosity with CSM interaction, Smith et al.\ (2007) proposed that the star suffered a tremendous explosive but non-terminal mass-loss event in the decade preceding the SN, analogous to the 19th century eruption of $\eta$~Carinae. Woosley et al.\ (2007) proposed a similar model, where the $\eta$~Car-like explosion in the decade preceding the SN could have been triggered by the pulsational pair instability (different from a genuine pair instability SN) in a star with initial mass 110 M$_{\odot}$. Interaction with an external CSM can, in principle, generate high luminosity as seen in other Type IIn SNe (e.g., Chugai et al.\ 2004). However, Smith et al.\ (2007) argued that this mechanism is difficult to reconcile with the relatively low progenitor mass-loss rate inferred from the weak X-rays detected by {\it Chandra} and the narrow H$\alpha$ emission in SN~2006gy's spectrum. In other words, signs of CSM interaction are seen in SN~2006gy, but that interaction appears far too weak to power the visual continuum. Even more troubling, the CSM interaction hypothesis cannot account for why the blast wave apparently did not decelerate while it was drained of more than 10$^{51}$ ergs (Smith et al.\ 2007). These inconsistencies, combined with the slow rise, low expansion speed, and high luminosity of SN~2006gy, led Smith et al.\ (2007) to appeal to the alternative of a pair-instability SN (Barkat et al.\ 1967; Bond et al.\ 1984; Heger \& Woosley 2002) where the core of the star is obliterated and the light is powered by radioactive decay of $\sim$10 M$_{\odot}$ of $^{56}$Ni. As we discuss below, however, it may be possible to account for the luminosity of SN~2006gy without placing such extreme demands on SN nucleosynthesis. The most pressing question about SN~2006gy boils down to this: Do we {\it need} a pair instability SN to provide a consistent explanation for its extremely high luminosity and spectral properties? The pair-instability hypothesis requires the introduction of an exotic phenomenon; while the idea has been around for decades, it has never been observed and is only expected to have occurred in the early Universe (e.g., Heger \& Woosley 2002). Published models of the light curves from pair instability SNe do indeed predict high luminosities, slow rise times, long durations, and slow expansion speeds (e.g., Scannapieco et al.\ 2005) qualitatively similar to SN~2006gy. The light curve shape depends on the assumed mass of the envelope, and most studies so far have been conducted for zero-metallicity stars with no mass loss. Further work is in progress, with various assumptions about the mass of the envelope (Nomoto et al.\ 2007; Kasen 2007; Young 2007). The progenitor of SN~2006gy is likely to have been a very massive star (Smith et al.\ 2007), but whether it was massive enough to suffer the pair instability remains an open question. Here we describe how a simple photon diffusion model can account for the high luminosity and long duration of SN~2006gy. Since the explosions of core collapse SNe typically deposit $E_0 \simeq 10^{51}$ ergs of energy in the SN envelope, an exceptional mass of $^{56}$Ni is not necessarily required for SN~2006gy if an efficient mechanism could convert the energy of the initial SN blast into emergent light. In SNe from compact progenitors, the light from the blast wave itself normally falls short by orders of magnitude. The trapped internal radiation will cool by virtue of adiabatic expansion, and its energy will consequently be tranferred to the kinetic energy of the expanding debris. The fraction of the blast energy emerging as radiation will be roughly $R_0/R_{max}$, where $R_0$ is the initial radius of the progenitor and $R_{max}$ is the radius of the SN photosphere at peak light. $R_0$ ranges from $\sim$10$^{-2}$ AU for the Wolf-Rayet progenitors of SNe Ib/c to $\sim$1 AU for the red supergiant progenitors of SNe~II, while typically, $R_{max} \simeq 100$ AU. Falk \& Arnett (1973, 1977) showed that a SN from a progenitor having a very extended envelope could produce a light curve with a long timescale and high luminosity, with very efficient conversion of blast wave energy into radiation. Below we consider this idea as a possible explanation for the light curve of SN2006gy. While we argue that this mechanism may provide a plausible explanation for the early light curve, we caution that it {\it does not negate} the possibility that SN~2006gy is powered by some other mechanism at late times, such as $^{56}$Co decay or continued CSM interaction. | The diffusion model for SN~2006gy that we present here does not require any new physical processes beyond those that account for core-collapse SNe. The only new ingredient is the hypothesis that the SN progenitor must have ejected an opaque shell having $\sim$10~M$_{\odot}$ extending to $\sim$160 AU shortly before it exploded. According to our model, the light curve of SN~2006gy observed up to about day 170 comes entirely from energy deposited by the initial blast and does not require any source of radioactive energy or ongoing CSM interaction. However, this internal energy source cannot last. As the shocked shell continues to expand, the radiative diffusion time becomes shorter than the time since explosion, after which the light curve will faithfully track the energy deposition rate. Most of the kinetic energy of the SN must reside in debris having velocity $\ga$4000 km s$^{-1}$, and thermal energy deposition from the impact of this debris with the expanding shell must terminate in $\la$1 yr. After that, there are only two other obvious sources of energy deposition to account for the light curve. The first is the decay of $^{56}$Co, which would produce an exponential tail to the light curve having luminosity $L \approx$1.4$\times$10$^{43}$ M($^{56}$Ni) $\exp(t/113.6 d)$ ergs s$^{-1}$, where M($^{56}$Ni) is the mass of newly synthesized $^{56}$Ni produced by the SN. If this is the explanation for the excess luminosity above the diffusion model after day 170, then about 8~M$_{\odot}$ of $^{56}$Ni would be required, indicated by the dashed line in Figure 1. On the other hand, failure to detect a continued exponential tail in the visual light curve at late times does not necessarily give an upper limit to M($^{56}$Ni) produced in SN~2006gy. The bolometric flux may shift to near-infrared wavelengths as ejecta cool, and it is quite possible that dust will form in the shocked shell, reprocessing the luminosity into far infrared radiation. This cooling may also cause changes in the bolometric correction as the SN evolves with time, and may play a role in the tail after day $\sim$160 seen in Figure 1. A second source could be continued shock interaction with the transparent CSM external to the opaque shell. We see clear signs of that interaction in the emission lines seen in the optical spectrum and in the faint X-ray emission (Smith et al.\ 2007). CSM interaction could produce a lower-luminosity tail to the light curve of SN~2006gy as long as the CSM density remains high. | 7 | 10 | 0710.3428 |
0710 | 0710.1331_arXiv.txt | We use high resolution (2048$^2$ zones) 2D hydrodynamic simulations to study the formation of spiral substructure in the gaseous disk of a galaxy. The obtained gaseous response is driven by a self-consistent non-axisymmetric potential obtained from an imposed spiral mass distribution. We highlight the importance of ultraharmonic resonances in generating these features. The temporal evolution of the system is followed with the parallel ZEUS-MP code, and we follow the steepening of perturbations induced by the spiral potential until large-scale shocks emerge. These shocks exhibit bifurcations that protrude from the gaseous arms and continue to steepen until new shocks are formed. When the contribution from the spiral potential relative to the axisymmetric background is increased from our default value, spurs protrude from the main arms after several revolutions of the gaseous disk. Such spurs overlap on top of the aforementioned shocks. These results support the hypothesis that a complicated gaseous response can coexist with an orderly spiral potential term, in the sense that the underlying background potential can be smooth yet drive a gaseous response that is far more spatially complex. | Young stars, H II regions and OB associations delineate the structure that gives its name to spiral galaxies (Elmegreen 1981; Vall\'ee 2002, 2005). This implies a correlation between the spiral structure and the process of star formation. The density wave theory attempts to explain the large-scale structure of these galaxies in terms of a wave propagating in the disk of stars. Yet this stellar wave alone cannot directly explain the narrow spiral arms as delineated by the products of star formation. Fujimoto (1968) first proposed that the dust lanes observed on the concave side (inside corotation) of spiral arms might be the result of a Galactic shock. Large scale shocks propagating in the gaseous layer of the Galaxy could induce a sudden compression of the interstellar medium (ISM) and trigger the star formation process. Since the temporal scales involved in this process are large (of the order of the Galactic rotation period), an isothermal approximation of the ISM is a first step in modeling the gaseous response, thus rendering a highly compressible-supersonic medium in which a small amplitude disturbance, such as that provided by an underlying spiral potential, may steepen into a shock wave. In a semi-analytical study of the gas flow in spiral density waves, Roberts (1969) found nonlinear steady-state solutions containing shocks. On the basis that these shocks were coincident with the imposed perturbing potential minima, he argued that the density jump could trigger star formation, and in this way narrow bands of young stars could delineate the spiral arms. Shu et al. (1973) studied the gas flow in the context of the spiral density wave by adopting a two phase model for the ISM. In their study they carried out a slightly nonlinear analysis of such flow and found that, at certain points in their solution, the amplitude became infinite, and they called these positions ultraharmonic resonances. They argued that additional prominent features could appear as a consequence of the ultraharmonic resonances, for instance: they found a secondary compression associated with the first one. However, these studies were carried out assuming steady state flow and tightly wound spirals. By removing the first constraint, Woodward (1975) showed how the convective steepening of the initial perturbation leads to large-scale shock formation (as a response to the driving potential) in a gaseous layer initially in circular orbit. He also found the effect of the first ultraharmonic resonance as a secondary peak in the density field, but numerical viscosity in his calculations inhibited the formation of secondary shocks. In two papers, Contopoulos \& Grosb\o l (1986, 1988) removed the second constraint and showed that nonlinear effects are not negligible in open spirals. Using numerical calculations, they demonstrated that for open spirals the effect of the first ultraharmonic resonance (which they call the 4:1 resonance) is such that stellar orbits do not support the spiral perturbation beyond the position of this resonance and hence the length of the stellar spiral arms is limited by this position. In an analytical study of the effects of the ultraharmonic resonances, Artymowicz \& Lubow (1992) explained this result as a cumulative effect of higher-order resonances between 4:1 and the corotation radius. They further argued that in the vicinity of the former the gas response is such that it looks like a 4-arm structure with a smaller pitch angle than that of the original two-arm pattern. It is clear that, due to the nature of the problem, the highly nonlinear phenomena associated with the gaseous response to a spiral density wave are best studied via numerical simulations. Results so obtained add interarm features and spiral substructure on top of large scale shocks. Patsis et al. (1997) found that in open spirals a bifurcation of the main spiral arms takes place at the 4:1 resonance position. By analyzing numerical experiments that include self-gravity in the study of the effects of the ultraharmonic resonances on gaseous disks, Chakrabarti et al. (2003, from here on CLS03) found secondary compressions, associated with the first ultraharmonic resonance, in models with a low Toomre parameter, Q, which were evolved only for a few revolutions of the disk. These compressions eventually became branches. In high-Q models, evolved for several revolutions they found the appearance of leading structures which they identified with spurs. However, spurs and branches not related to resonances have also been obtained in numerical simulations (Dobbs \& Bonnell 2006; Wada \& Koda 2004). By taking into account frozen magnetic fields and self-gravity in their local arm simulations, Kim \& Ostriker (2002, 2006) showed that gaseous spurs form as a consequence of gravitational instability inside the spiral arms, a result confirmed with global 2D simulations by Shetty \& Ostriker (2006). In those simulations spurs jut out of the spiral arms at regular intervals. In this paper we carry out two dimensional, global, high resolution hydrodynamic simulations to further study the formation of spiral substructure in galactic gaseous disks. Our approach differs from previous work in that we employ a self-consistent model for the spiral stellar potential (in the orbital sense). This potential accounts for its own self-gravity and thus is no longer a local arm approximation to the driving term, making it more appropriate for simulations of open galaxies. This paper is organized as follows: in \S 2 we describe the potential we employ, in \S 3 we present the obtained gaseous response and the type of substructure related to it, in \S 4 we discuss the results and compare with previous work, and in \S 5 we present our conclusions. | A fully nonlinear treatment of the propagation of spiral density waves in a gaseous disk is tractable using numerical simulations. By considering an open spiral, in the absence of self-gravity in the gaseous layer, we are able to obtain rich substructure associated with an external spiral potential. Our work differs from other published results in that we employ a self-consistent driving term that considers its own gravity and removes the local arm approximation. We have presented experiments showing the formation of a four-arm structure in response to a two-arm driving pattern. Initially the gas accumulates in the potential minima. After some time the gas responds to resonances and the main arms bifurcate. These features have a considerable azimuthal extension and continue to steepen with time. Eventually a four-arm shock structure emerges, in the gaseous layer that coexist with the underlying two-arm potential, made up by the stars. This result appears to agree with observational data published by Drimmel (2000) where he concludes that, using optical tracers, the spiral structure in our Galaxy is best fitted by a four-arm structure, while in the infrared a two-arm structure dominates. If we place the first ultraharmonic resonance outside the region of influence of the spiral (in our case by changing the value of $\Omega_p$) the main arms do not bifurcate and large-scale shocks are the only induced response in the gaseous layer. In this way we show that if the first ultraharmonic resonance is placed outside the region of influence of the perturbing term, no spiral substructure appears, thus emphasizing the connection between this resonance and the bifurcations of the gaseous arms. If we combine the effects of the resonances with a large forcing amplitude we obtain, on top of this four-arm structure, spurs protruding from the main arms in a region between the inner Lindblad resonance and the first ultraharmonic resonance. Overlapping of nonlinear effects take place in our long-timescale global simulations, showing that an orderly spiral density wave potential can produce a gaseous response that is strongly disordered. The inclusion of additional processes such as magnetic fields, self-gravity for the gaseous layer, and thermal processes in the ISM may lead to the appearance of additional substructure and their consequent fragmentation into bound condensations. Numerical experiments addressing these topics are needed in order to improve our understanding of the global ISM in galaxies and its relation to large-scale processes. | 7 | 10 | 0710.1331 |
0710 | 0710.1277_arXiv.txt | We present new deep {\it Chandra} observations of the Centaurus A jet, with a combined on-source exposure time of 719 ks. These data allow detailed X-ray spectral measurements to be made along the jet out to its disappearance at 4.5 kpc from the nucleus. We distinguish several regimes of high-energy particle acceleration: while the inner part of the jet is dominated by knots and has properties consistent with local particle acceleration at shocks, the particle acceleration in the outer 3.4 kpc of the jet is likely to be dominated by an unknown distributed acceleration mechanism. In addition to several compact counterjet features we detect probable extended emission from a counterjet out to 2.0 kpc from the nucleus, and argue that this implies that the diffuse acceleration process operates in the counterjet as well. A preliminary search for X-ray variability finds no jet knots with dramatic flux density variations, unlike the situation seen in M87. | \label{intro} Low-power, Fanaroff-Riley class I (FRI) \citep{fr74} radio galaxies are numerically the dominant population of radio-loud active galaxies in the universe. Their dynamics and energy content are thus essential to models of AGN `feedback' in which the energy released in the process of accretion onto the central supermassive black hole is transported to large spatial scales via the interaction between the jets and the external medium. To understand the dynamics and energetics of these sources it is crucial that we should understand the particle acceleration processes by which the bulk kinetic energy of the jets is translated into the internal energy of relativistic plasma. The discovery with {\it Chandra} that X-ray emission is common in the inner few kpc of FRI jets \citep{wbh01} provides a very strong argument that {\it in situ} high-energy particle acceleration is taking place in those regions. The broad-band spectral energy distribution and X-ray spectrum of the X-ray emission imply a synchrotron origin for the X-rays \citep{hbw01}. For magnetic fields close to the equipartition value in a typical powerful jet, the loss timescale ($\tau = -E/\dot E$) for the $\gamma \sim 10^7$ -- $10^8$ electrons emitting X-ray synchrotron emission is tens of years. Thus observations of X-ray synchrotron emission essentially tell us where particle acceleration is happening {\it now}: in particular, resolved, diffuse X-ray emission implies a particle acceleration mechanism that must be distributed throughout the jet. To probe the nature of the acceleration mechanisms in FRI jets we require observations that can reliably distinguish between compact and diffuse X-ray emission: this implies (at the most optimistic) a spatial resolution that is comparable to the loss spatial scale, the distance travelled by an electron before synchrotron losses remove it from the X-ray band, or roughly $c\tau$. Even with the comparatively high angular resolution of {\it Chandra}, this can only be achieved in the nearest FRI radio galaxy, Centaurus A. Given Cen A's distance of $\sim 3.7$ Mpc (the average of 5 distance estimates to Cen A: \citealt{fmst07}), {\it Chandra}'s resolution corresponds to $\sim 10$ pc. {\it Chandra} observations of Cen A show a complex, knotty and in places edge-brightened jet structure (\citealt{kfjk00, kfjm02, hwkf03} [hereafter H03]; \citealt{ksat06} [K06]; \citealt{hkw06} [H06]). It has been argued (H03,K06) that at least two acceleration mechanisms are required: the dynamics and spectra of the compact knots in the inner part of the jet are consistent with a shock origin, while the diffuse emission is inconsistent both spectrally (H03) and in terms of the number distribution of knots (K06) with being the sum of many unresolved knots with the same properties as those observed, and is instead probably truly diffuse, implying a distributed acceleration mechanism such as second-order Fermi acceleration \citep[e.g.,][]{so02} or magnetic field line reconnection \citep[e.g.,][]{bl00}. Here we present the results of new, much deeper {\it Chandra} observations and their consequences for the nature of particle acceleration in the Cen A jet and counterjet. \begin{figure*} \epsfxsize 16.6cm \epsfbox{f1n.eps} \caption{False-color image of the jet and counterjet regions of Cen A using all ACIS observations. The {\it Chandra} data from each observation have been exposure corrected at an appropriate energy and then combined, weighting according to the value of the exposure map. The images are binned in standard {\it Chandra} pixels ($0\farcs492$ on a side) and smoothed with a FWHM = $1\farcs0$ Gaussian. Red shows exposure-corrected counts in the energy range 0.4--0.85 keV, green shows 0.85--1.3 keV and blue 1.3--2.5 keV. Contours [$7 \times (1,4,16\dots)$ mJy beam$^{-1}$] are from the 5-GHz VLA map of H06 with $6\farcs0$ resolution.} \label{fc} \end{figure*} In this analysis {\it Chandra} data processing was done using {\sc ciao} 3.4 and {\sc caldb} 3.3.0.1. Spectral fitting was carried out in {\sc xspec}. We define the spectral index $\alpha$ such that flux $\propto \nu^{-\alpha}$: the photon index $\Gamma = 1+\alpha$. Errors quoted and plotted throughout are $1\sigma$ for one interesting parameter. | The inner part of the Cen A jet is dominated by the shock-related knots discussed by H03, although we have so far failed to see the dramatic variability seen in M87, interpreted as large changes in particle acceleration processes or in jet flow. The radio-associated compact features in the counterjet are likely also the result of small-scale shocks, presumably caused as the counterjet flow encounters compact obstacles. However, the outer 3.4 kpc of the jet -- and possibly also the newly detected extended X-ray counterjet components -- are different from the shock-dominated regions in their X-ray spectra and structure and in their radio/X-ray ratio. Our results add to the evidence that an unknown, distributed particle acceleration process operates in the jet of Cen A. Diffuse X-ray emission with a similar broad-band spectrum is seen in the jets of many other FRI sources (e.g., 3C\,66B: \citealt{hbw01}) and it also resembles the extended synchrotron X-ray emission seen in the jets (e.g., K06) and hotspots \citep[e.g.,][]{hck07} of the more powerful FRIIs. The Cen A data allow us to study the spatial and spectral properties of this emission and the physical processes responsible for it in more detail than is possible in any other object. \vspace{-10pt} | 7 | 10 | 0710.1277 |
0710 | 0710.2332_arXiv.txt | \baselineskip=13pt A new mass table calculated by the relativistic mean field approach with the state-dependent BCS method for the pairing correlation is applied for the first time to study $r$-process nucleosynthesis. The solar $r$-process abundance is well reproduced within a waiting-point approximation approach. Using an exponential fitting procedure to find the required astrophysical conditions, the influence of mass uncertainty is investigated. $R$-process calculations using the FRDM, ETFSI-Q and HFB-13 mass tables have been used for that purpose. It is found that the nuclear physical uncertainty can significantly influence the deduced astrophysical conditions for the $r$-process site. In addition, the influence of the shell closure and shape transition have been examined in detail in the $r$-process simulations. | It is of the utmost interest to explore the ``terra incognita'' of exotic nuclei, as evidenced by the fact that several Radioactive Ion Beam (RIB) facilities are being upgraded, under construction or planned to be constructed worldwide. Such investigations of the properties of these exotic nuclei, which may behave very differently from the nuclei around the $\beta$-stability line, result in new discoveries such as the halo phenomenon~\cite{THH-PRL85,Schwab94} - nucleons spread like a thin mist around the nucleus, which can significantly increase the nuclear reaction ratio. Stellar nucleosynthesis processes such as the $r$-process~\cite{Cowan-PR91,Qian-PPNP03}, which is responsible for roughly half of the enrichment of elements heavier than iron in the universe, also require a thorough understanding of the properties of exotic nuclei. Key properties like masses for example, determine the path that the nucleosynthesis process follow in the nuclei chart. Nevertheless, despite many experimental efforts, present knowledge of exotic nuclei still does not include much of what is required for a complete understanding of $r$-process nucleosynthesis. After the first systematic introduction to the $r$-process~\cite{BBFH} half a century ago, $r$-process calculations for a long time could only rely on the phenomenological nuclear droplet mass formula~\cite{Hilf} because of the lag of both experimental and theoretical development. Fortunately, in the last 15 year the theoretical study of nuclear properties has made tremendous progress and $r$-process calculations~\cite{Kartz-APJ93,pfeiffer-ZPA97,Wanajo-APJ04} have been carried out based on the refined droplet model FRDM~\cite{FRDM}, Hartree-Fock approach like ETFSI-Q~\cite{ETFSIQ}, and the very recent microscopic rooted Hartree-Fock Bogliubov (HFB)~\cite{HFB,HFB2,Samyn-PRC03}. Despite the progress in the theoretical nuclear structure physics, mass models predictions (which by design concentrate on different nuclear structure aspects) still show a large deviation when going to very neutron-rich nuclides, even though they have achieved similar quality to describe known nuclides. This is specially troublesome since the astrophysical scenario in which an $r$-process may occur is a matter of debate and all astrophysical simulations dedicated to the nature of the stellar environment depend on the input from nuclear physics. Mass model predictions, even in models that give similar global $rms$ error still show local deviations differently. In principle, microscopic-rooted mass models should have a more reliable extrapolation to the unknown regions, therefore these studies have received more and more interest as evidenced by the increasing number of non-relativistic HFB investigations~\cite{HFB,HFB2,Samyn-PRC03,HFB-13}. Based on a mass-driven fitted method, the latest HFB models have achieved a similar quality ($rms\sim$ 0.7 MeV) as the phenomenological FRDM mass model for known masses. More recently, another microscopic-rooted approach, the relativistic mean-field (RMF) theory~\cite{Wale-AP74} has received broad attention due to its successful description of several nuclear phenomena during the past years (for recent reviews, refer to Refs.~\cite{Bender-RMP03,Meng06}). In the framework of the RMF theory, the nucleons interact via the exchanges of mesons and photons. The corresponding large scalar and vector fields, of the order of a few hundred MeV, provide simple and efficient descriptions of several important phenomena such as the large spin-orbit splitting, the density dependence of optical potential, the observation of approximate pseudo-spin symmetry, etc.. Moreover, the RMF theory can reproduce well the isotopic shifts in the Pb region~\cite{SLR-PLB93}, explain naturally the origin of the pseudo-spin symmetry~\cite{AHS-PLB69,HA-69} as a relativistic symmetry~\cite{Ginocchio-97,MSYP-PRC98,MSYA-PRC99,Chen-CPL03} and spin symmetry in the anti-nucleon spectrum~\cite{ZMP-PRL03}. The first RMF mass table was reported in Ref.~\cite{Hirata-NPA97} for 2174 even-even nuclei with $8\leq Z \leq 120$ but without including pairing correlations. Later on, the calculation was improved by adopting a constant-gap BCS method and calculated 1200 even-even nuclei with $10 \leq Z \leq 98$~\cite{Lalazissis-ADNDT99}, most of which are close to the $\beta$-stability line. More recently, using the state-dependent BCS method with a $\delta$-force~\cite{Sanulescu-PRC00,Geng-PTP03}, the first systematic study of the ground state properties of over 7000 nuclei ranging from the proton drip line to the neutron drip line was performed \cite{RMFBCS}. Comparison of this calculation with experimental data and to the predictions of other mass models will be presented in more detail in Sec II. Considering the recent development of the microscopic mass models in both the HFB and RMF approach, it is very interesting to examine their applicability to an $r$-process calculation. The main goals of this paper is to explore to what extent the solar $r$-process abundance can be reproduced using the new RMF mass table~\cite{RMFBCS} and by comparing with other theoretical mass models, to determine the influence of nuclear mass uncertainty in $r$-process calculations. The paper is organized as follows. In Sec.~\ref{sec:RMF} the global agreement of the new RMF mass table with the experimental data is discussed and the RMF prediction in the very neutron-rich range is compared with the FRDM~\cite{FRDM}, the ETFSI-Q~\cite{ETFSIQ} and the latest HFB-13~\cite{HFB-13} mass tables. In Sec.~\ref{sec:$r$-process}, a short introduction to a site-independent $r$-process approach is given. In Sec.~\ref{sec:calculation}, the new mass table is applied to reproduce the solar $r$-process abundances. In addition, the result is compared to the $r$-process abundances obtained with the predictions of the FRDM, ETFSI-Q and HFB-13 mass models. Finally the summary and conclusions are given in Sec.~\ref{sec:conclusion}. | Summary} We have applied the most recent comprehensive mass models, the non-relativistic microscopic-rooted HFB-13 and the relativistic RMF in $r$-process calculations. For the sake of comparison, we also included the widely used macro-microscopic models FRDM and ETFSI-Q. Of these models, the HFB-13 and RMF models are used for the first time in such calculations. Based on a simple $r$-process model, it is found that all mass models reproduce the main features of the solar $r$-process pattern and the position of the abundance peaks. Since $r$-process simulations have to rely on predicted nuclear physics properties of unknown regions in the nuclear chart, we have compared the predictions of different mass models. We have also made a systematic study of the influence due to the mass model uncertainty in the application of the $r$-process and thus in the required astrophysical conditions. This nuclear physical uncertainty is very important for the complete understanding of the $r$-process since the results of more modern full dynamic $r$-process calculations depend on the nuclear mass input used. It is found that the deduced astrophysical conditions like the neutron irradiation time of the $r$-process can be significantly different depending on the mass model used. Among the different models, the simulation using the RMF masses requires a longer time scale (up to a factor of 4) than those using FRDM and HFB-13 models. Furthermore, it is found that the optimal astrophysical conditions obtained using the ETFSI-Q and RMF mass models require a relatively constant weighting factor for neutron densities in the range $10^{22}$ to $10^{28}$ cm$^{-3}$, while the FRDM and HFB-13 simulations favor a large weighting factor at low densities. In addition, we have explored the possible deficiencies in different mass models, and found that the observed abundance underproduction before the abundance peaks in all the models can be a combined and complex effect of both shell structure and shape transition. An exception is the underproduction at $A\sim 115$ in the HFB-13 model which can be attributed to incorrect $\beta$-decay rates. Future experiments are needed to determine the strength of the shell closure towards the neutron drip line as well as the precise locations of the shape transition toward the shell-closures. | 7 | 10 | 0710.2332 |
0710 | 0710.2104_arXiv.txt | In this work we show that 3-3-1 model with right-handed neutrinos has a natural weakly interacting massive particle (WIMP) dark mater candidate. It is a complex scalar with mass of order of some hundreds of GeV which carries two units of lepton number, a scalar bilepton. This makes it a very peculiar WIMP, very distinct from Supersymmetric or Extra-dimension candidates. Besides, although we have to make some reasonable assumptions concerning the several parameters in the model, no fine tunning is required in order to get the correct dark matter abundance. We also analyze the prospects for WIMP direct detection by considering recent and projected sensitivities for WIMP-nucleon elastic cross section from CDMS and XENON Collaborations. | \setcounter{equation}{0} \label{sec1} The problem of matter density in the Universe seems to be one of the most intriguing and exciting subjects in modern Physics. The growing refinement achieved in cosmological data leaves no doubt about a dark component in the observed mass density, constituting roughly $22\%$ of all energy density acording to the three year run of WMAP~\cite{WMAP3}. This yet unknown component has to be non-baryonic, its interaction with the electroweak Standard Model (SM) particles should be negligible and it has to be cold, i.e., non-relativistic at the time it decouples from the radiation bath, the so called Cold Dark Matter (CDM). From the theoretical side, there are some proposals to explain the CDM in the context of Particle Physics models (see Ref.~\cite{DMmodels,Murayama,dmoutros,Dobrescu,little} and references therein for a review of the subject). Among them there are models which present natural candidates to play this role, the weakly interacting massive particles (WIMP)'s, with mass ranging from approximately $1$~GeV to $1$~TeV. These WIMP's are nice candidates because their masses are in the GeV realm, turning them cold at decoupling, and mainly because their weakly interacting aspect not only yields a thermally averaged annihilation cross section of order of weak interactions, leading to the expected order of magnitude to CDM abundance, but also coincides with the scale of Particle Physics models to be probed at the Large Hadron Collider (LHC) that finds itself at the final stage to start its running phase~\cite{LHC}. It also presents the possibility of being seen in direct detection experiments since its massiveness would imply an observable recoil of nuclei in elastic collisions~\cite{DMmodels,wimpexperiments}. The most promising scenarios where such WIMP's can be present in the particle spectrum are Supersymmetry (SUSY) and Extra Dimensions models~\cite{DMmodels,Murayama,Dobrescu}. All these models dispose of some kind of discrete symmetry in order to stabilize their CDM candidates. Also, they have to be realized at the electroweak scale so that their new particles are potentially good candidates for CDM. Although such models may represent the greatest expectations in Particle Physics for the Physics at TeV scale to be probed by LHC, the absence of any experimental evidence to support these models allows us to work with alternatives. One of such alternative routes concerns the enlargement of the gauge symmetry group from $SU_C(3)\otimes SU_L(2)\otimes U_Y(1)$ to larger groups. In particular, there exists a simple extension of the SM gauge group to $SU_C(3)\otimes SU_L(3)\otimes U_X(1)$, the so called 3-3-1 model~\cite{PleitezPisano}. This class of models is interesting for several reasons, not only because they mean a different scenario, but because they possess several nice features. For example, (i) the family problem is absent in this model since it demands that there be only three families of fermions when anomalies are canceled and asymptotic freedom is considered~\cite{PleitezPisano}; (ii) electric charge quantization is automatic~\cite{qcharge}; (iii) right-handed neutrinos can be part of the spectrum in some versions of the model~\cite{331valle,331RH,singer} and their tiny observed mass difference can be easily accommodated~\cite{lightnu}; (iv) axions and majorons are a natural outcome in some versions~\cite{axionmajoron}, providing light particles which could also contribute to the problem of dark matter origin. Besides, it is possible that a custodial symmetry exists in these models which would make them indistinguishable from SM at low energy scales~\cite{331np}, and it would then be a strong rival to SM itself. In this work we will concentrate on the 3-3-1 version of the model with right handed neutrinos in the spectrum (3-3-1$RH\nu$) model~\cite{331RH}. The reason behind this choice relies on the fact that neutrinos mass is already a mandatory property that needs to be included in all reasonable extensions of SM. Besides, the model can be implemented with just three scalar triplets instead of including a sextet as in other versions, being economical in its content. However, the most appealing motivation to deal with 3-3-1$RH\nu$ to explain the origin of CDM is due to the possibility of having a candidate with a very distinct signature. Among its properties the model can be made lepton number conserving if some of its fields carry two units of lepton number which will be called bileptons. This peculiar property has many phenomenological implications. Namely, rare lepton decays can emerge, neutrinoless double beta decay is allowed, right handed neutrinos are going to appear as byproducts of heavy vector bileptons decays and so on. It is then automatic to ask if some of these additional bilepton fields can be a CDM candidate, once it is provided with a very specific quantum number appropriate to forbid its interaction with many of the electroweak fields. This would play a similar role as that played by the discrete symmetries in the competing models cited above. Moreover, as the sought candidate is merged in the exotic new effects just mentioned, their appearance in the coming collider experiments would represent an unquestionable evidence of our CDM candidate. What we are going to investigate in this work is the possible realization of this scenario in the 3-3-1$RH\nu$ model for one of the bilepton scalars. We are interested in identifying this field and characterize it as a WIMP. This can be realized if it can be shown that it is stable in the range of parameters for which the abundance is in accordance with the recent data of WMAP~\cite{WMAP3}. It is important to stress that the natural perturbative scale for the 3-3-1 models is on the TeV scale~\cite{alexlimit}, which is suitable for obtaining a WIMP. In view of all this, it seems that such a possibility is as welcome as any previous attempt to explain the CDM content through SUSY or Extra dimensions, offering a completely new and distinct particle to do this job. It should be mentioned that other works exist in the literature trying to explain CDM in the context of 3-3-1 models~\cite{DM331}, but their aim was to obtain a self-interacting dark matter to avoid excessively dense cores in the center of galaxies and clusters as well as excessive large number of halos within the local group when contrasted to observations~\cite{spergel}, which demands a light dark matter candidate. We do not pursue this approach here. On the contrary, we want to show that the 3-3-1$RH\nu$ model possesses a bilepton scalar that can play the role of a WIMP, which is the preferred candidate for CDM. Besides, we get this without the need to fine tune its couplings to small and unnatural values. This work is organized as follows. In Sec.~\ref{sec2} we present the model with its content and interactions. In Sec.~\ref{sec3} we diagonalize the mass matrices for the particle spectrum, allowing us to characterize the new extra particles. In Sec.~\ref{sec4} we identify our WIMP with one of the neutral scalar bileptons of the model, checking its viability as CDM candidate by computing its abundance for a range of values of the parameters, which turns out to be natural in the sense we do not need to make any fine adjustment on the parameters. Then we study the prospects of direct detection for this WIMP. We conclude with Sec.~\ref{sec5}. | \label{sec5} The 3-3-1RH$\nu$ model admits a couple of bilepton particles in its spectrum, raising the possibility of having a CDM candidate, since bileptons carry two units of lepton number. This is so because such a very specific quantum number is appropriate to forbid its interaction with many of the electroweak fields, since they are allowed to decay only on other bileptons, which have to be heavier than SM particles. Considering this scenario we obtained the particle mass spectrum of scalars in 3-3-1RH$\nu$ model and, by assuming some conditions over the parameter space, we have shown that the lightest bilepton in the model turns out to be a scalar, a combination of two scalar interaction eigenstates that we called $\phi$. For the region where the values of the parameters guarantee the stability of this scalar, we computed the $\phi$ abundance and obtained stringent constraints for the parameters in order to have agreement with WMAP results for CDM abundance. We have found that $\phi$ can have mass ranging from about $600$~GeV to some Few TeV, characterizing it as a heavy WIMP. It is opportune to say that, although we have restricted our parameter space due to lack of knowledge on several couplings in the model, we had no need to unnaturally adjust them to very small values as generally happens in several models, including Supersymmetry. In fact we assumed that these couplings are close to one, and checked that $\phi$ is an excellent candidate to represent a WIMP and explain the presently observed CDM abundance. We also studied the possibility of observing this WIMP in direct detection experiments. For this we have computed the elastic scattering $\phi$-nucleon cross section and contrasted our results with present and future experiments. We have seen that $\phi$ is still far from the range of detection for current and near future CDMS and XENON sensitivities, at least for the short parameter space considered in this work. However, even this limited scenario can be at reach for projected Phases B and C of Super-CDMS~\cite{SuperCDMS}. Besides this, it would be interesting to pursue the production of $\phi$ at collider experiments, mainly at LHC, and also extend our search including a larger region of the parameter space considering additional phenomenological constraints on 3-3-1RH$\nu$ model from Collider physics and Cosmology, as well as include prospects for $\phi$ indirect detection too, a gap we wish to fill soon. Finally, we would like to stress that our proposed WIMP is not only feasible but a reasonable alternative in the sense that the Particle Physics model we are dealing with is only a small extension of the SM gauge group, whose scale is about to be assessed at LHC. There are several features that distinguishes the 3-3-1RH$\nu$ model from other extensions, like SUSY and Extra dimensions models. Namely, we have not only bilepton scalars in the spectrum but vector bosons and quark bileptons, all of them acquiring mass at hundreds of GeV. Certainly their signal at detectors are worth to be studied. Besides, new phenomena are predicted in this model~\cite{331valle,lightnu,axionmajoron,SCPV}, including neutrinoless double beta decay, rare decays, new sources of CP violation and so on. Presence of such signals would reinforce our expectation concerning a bilepton WIMP to explain CDM in the Universe. \\ \\ \noindent {\bf Acknowledgments:}\\ The authors acknowledge the support of the Conselho Nacional de Desenvolvimento Cient\'{\i}fico e Tecnol\'ogico (CNPq). | 7 | 10 | 0710.2104 |
0710 | 0710.1727_arXiv.txt | Stellar kinematics show no evidence of hidden mass concentrations at the centre of M83. We show the clearest evidence yet of an age gradient along the starburst arc and interpret the arc to have formed from orbital motion away from a starforming region in the dust lane. | \label{sec:intro} The nucleus of M83 is offset from the bulge centre and surrounded by a semicircular starburst arc (Fig. 1). \citet[][hereafter H01]{Harris01} dated the star clusters in the arc with WFPC2 photometry and found evidence of an age gradients along it. However, the reddening vector parallelled the tracks in the two-colour diagram and clusters may be overlooked or poorly sampled in the visible because of extinction. \citet[hereafter T00]{Thatte00} proposed the existence of a second mass concentration after NIR long-slit kinematics revealed an additional peak in the velocity dispersion, 2.7\hbox{$^{\prime\prime}$}\ SW of the nucleus. Further studies with IFUs linked velocity gradients in gas kinematics to hidden mass concentrations at different locations: \citet{Mast06} report a gradient in \textrm{H}{\large $\alpha$}\ at 3\hbox{$.\!\!^{\prime\prime}$}9$\pm$0\hbox{$.\!\!^{\prime\prime}$}5 W of the nucleus while Diaz et al. (\citeyear[hereafter D06a and D06b]{D06a,D06b}) report a gradient in \textrm{Pa}{$\beta$}\ 7\hbox{$^{\prime\prime}$}\ WNW of the nucleus. However, gas kinematics are known to suffer non-gravitational effects \citep{KR95}. We have analysed new VLT ISAAC K band long-slit spectroscopy together with archival HST data (\textrm{Pa}{\large $\alpha$}\ from NICMOS and \textrm{H}{\large $\alpha$}\ from WFPC2). Fig. 1 illustrates slit positions on the HST data. The combined data gives equivalent width (EW) measurements of \textrm{H}{\large $\alpha$}, \textrm{Pa}{\large $\alpha$}\ and CO (2.3\hbox{$\umu$m}), as well as stellar and gas kinematics. \begin{figure} \centering \begin{minipage}[t]{0.7\textwidth} \centering \resizebox{0.9\textwidth}{!}{\includegraphics{HOUGHTON_fig1.ps} } \end{minipage}% \begin{minipage}[c]{0.3\textwidth} \vspace{-1.5cm} \caption[]{\textbf{(a)} A BVR image of M83 (logarithmic intensity) using data from \citet{Larsen99}, with the footprint of the homogenised HST data overlaid. Axes are scaled in degrees. \textbf{(b)} F222N NICMOS image of the central 20\hbox{$^{\prime\prime}$}$\times$20\hbox{$^{\prime\prime}$} with ISAAC slit positions overlaid. Positions of the putative hidden mass concentrations are also shown as purple (T00) and red (D06a,D06b) triangles. Flux is given in ergs s$^{-1}$ cm$^{-2}$ and axes are scaled in arcseconds. } \end{minipage} \label{fig:m83} \end{figure} \begin{figure} \centering \begin{minipage}[r]{0.6\textwidth} \centering \vspace{-2.5cm} \includegraphics[width=0.7\textwidth]{HOUGHTON_fig2.ps} \end{minipage}% \begin{minipage}[c]{0.4\textwidth} \caption[]{An image of the centre of M83 with ages overplotted; note that the circumnuclear arc extends into and along the dust lane; the position, scaling and orientation is the same as Fig. 1b; NICMOS F222M is red, WFPC2 F814W is green and WFPC2 F300W is blue.} \end{minipage} \label{fig:sb99} \end{figure} | 7 | 10 | 0710.1727 |
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0710 | 0710.1382_arXiv.txt | Planned space-based ultra-high-energy cosmic-ray detectors (TUS, JEM-EUSO and S-EUSO) are best suited for searches of global anisotropies in the distribution of arrival directions of cosmic-ray particles because they will be able to observe the full sky with a single instrument. We calculate quantitatively the strength of anisotropies associated with two models of the origin of the highest-energy particles: the extragalactic model (sources follow the distribution of galaxies in the Universe) and the superheavy dark-matter model (sources follow the distribution of dark matter in the Galactic halo). Based on the expected exposure of the experiments, we estimate the optimal strategy for efficient search of these effects. | \label{sec:intro} The next step in the studies ultra-high-energy (UHE) cosmic rays (CRs) is related to the use of space-based detectors observing the fluorescent light induced by air showers in the terrestrial atmosphere with large exposures. It is expected that these instruments will help us to shed light on the origin of the highest-energy cosmic particles which remains unknown up to now. One of the important signatures of particular UHECR models is the global anisotropy of arrival directions of the highest-energy events. For instance, models where the origin of UHECRs is attributed to the acceleration in astrophysical objects (so-called ``bottom-up'' scenarios) naturally predict that the distribution of arrival directions follows the distribution of these cosmic accelerators. In the most common scenario with extragalactic protons and/or nuclei, the patterns of the distribution of galaxies in the nearby Universe should be seen~\cite{EGmogila} on the UHECR skymap because of the limited propagation length of these particles due to the GZK effect~\cite{GZK} or nuclear photodisintegration. Strong suppression of the cosmic-ray flux is predicted at highest energies in these models. On the other hand, the ``top-down'' models~\cite{TD} (see Refs.~\cite{TDrev} for reviews and references) with the distribution of sources in the Galactic halo following that of the dark matter (which is the case for the superheavy dark-matter (SHDM) particles and some of topological defects) predict~\cite{DubTin} the Galactic center-anticenter asymmetry due to the non-central position of the Sun in the Galaxy (see Refs.~\cite{DMmogila,ABK1,ABK2} for extensive discussions). Models of this kind predict continuation of the cosmic-ray spectrum and gamma-ray dominance beyond $10^{20}$~eV. Currently, neither the spectrum nor anisotropy observations can definitely favour one of the two scenarios at the highest energies, $E \sim 10^{20}$~eV. Indeed, the AGASA experiment claims~\cite{AGASAspectrum} the super-GZK continuation of the spectrum while the HiRes collaboration reports~\cite{HiRes-cutoff} the observation of the Greisen--Zatsepin--Kuzmin (GZK) cutoff (data of other experiments, including recent results of the Pierre Auger Observatory (PAO)~\cite{PAO-SDspectrum,Engel-CIC}, are not yet conclusive: though the unsuppressed continuation of the spectrum is excluded by the Auger data, the cutoff in the spectrum is not clearly seen). The limits on the gamma-ray fraction in the primary cosmic-ray flux (currently the most restrictive ones arise from the AGASA and Yakutsk muon data at $E>10^{20}$~eV~\cite{gamma}, from the Yakutsk muon data at $E>4 \cdot 10^{19}$~eV~\cite{Yak-gamma} and from the Pierre Auger Observatory data on the shower geometry at $E>2 \cdot 10^{19}$~eV and $E>10^{19}$~eV~\cite{PAO-gamma}) disfavour the SHDM scenario and even exclude it for particular values of the dark-matter parameters (see e.g.\ Ref.~\cite{Risse} for a more detailed discussion of some of these limits). Current experiments do not report any significant deviations from the global isotropy at the highest energies\footnote{ After this paper was submitted, PAO reported a significant deviation from isotropy~\cite{PAOagn}. The definite interpretation of this effect awaits further data (see e.g.\ Ref.~\cite{Comment} for discussion). }. This is however not conclusive both due to the low statistics and due to a limited field of view of any terrestrial-based installation. The steeply falling spectrum of cosmic rays makes it very difficult to obtain a reliable measure of the global anisotropy in {\em any} combination of terrestrial experiments. Indeed, the relative systematic difference in the energy estimation between two installations located in different parts of the Earth (and thus observing different parts of the sky) can hardly be made smaller than some $\sim 15\%$. Such a relative error would give $\sim 30\%$ higher integral flux seen by one of the experiments with respect to the other at the same reconstructed energy. Thus, possible observations of global anisotropy can be attributed both to a physical effect and to unknown systematics in the energy determination. Moreover, a seemingly isotropic distribution of the arrival directions over the sky might represent a physically anisotropic one masked by the systematic effects. On the other hand, the planned space-based experiments, e.g.\ TUS~\cite{TUS}, JEM-EUSO~\cite{EUSO} or S-EUSO~\cite{S-EUSO}, will provide a unique opportunity to observe full sky with a single detector. While not being free from systematic uncertainties in the energy determination, an experiment of this kind would not introduce strong direction-dependent systematics and thus would be able to perform the studies of the global anisotropy at high confidence. One may expect that in future space-based experiments with their huge exposures, particular sources of UHECR will be determined on event-by-event basis for the case of the astrophysical scenario. This is not an easy task, however: limited angular resolution together with large numbers of events would lead to identification problems similar to thoose of the gamma-ray astronomy\footnote{See e.g.\ Refs.~\cite{EGRET-statistics} for discussions of statistical methods of analysis of photon-by-photon EGRET data. Currently, the gamma-ray sky at energies $(\sim 0.1 \div 1)$~GeV contains 101 identified source, 170 unidentified ones and a strong non-uniform diffuse background of unidentified origin.} but strongly enhanced due to magnetic deflections of the charged cosmic-ray particles. The searches for global patterns in the distribution of UHECR arrival directions will thus be important in any case. A robust method for the study of any global asymmetry in the arrival directions is the harmonic analysis (see e.g.\ Ref.~\cite{Sommers}). It works perfectly if the predicted effect may be clearly seen in the first few harmonics but requires large statistics to reveal/exclude more complicated patterns. Here, we determine the optimal strategy for the searches of global anisotropies even with the low-resolution experiments (TUS) and for short exposure times. The strategy is to fix two regions of the sky (not necessary covering full $4\pi$) which are expected to provide the strongest contrast in over/underdensity of events with respect to the null hypothesis of the isotropic distribution. The shape and the size of these regions, as well as the energy range of the events, are determined {\it a priory} in order to balance the strength of the expected anisotropy (increasing for smaller regions and higher energies) and its statistical significance. The aim of this paper is to simulate optimal regions and energies for TUS, JEM-EUSO and S-EUSO, suitable to test the scenarios of extragalactic sources and of decays of superheavy dark matter concentrated in the Galactic halo. To this end, we perform new and improved (with respect to previous works) simulations of the distribution of arrival directions expected in both cases. The paper is organised as follows. In Sec.~\ref{sec:EG} and Sec.~\ref{sec:DM}, we discuss our simulations for extragalactic sources and for SHDM decays, respectively. Sec.~\ref{sec:experiments} gives specific predictions for three coming spaceborn experiments, TUS, JEM-EUSO and S-EUSO. Sec.~\ref{sec:concl} contains our brief conclusions. Some technical details are collected in \ref{app:density-function}. | \label{sec:concl} In this work, we made quantitative predictions for the global anisotropy of the UHECR arrival directions expected in two distinct scenarios (``top-down'' and ``bottom-up'') of the origin of the highest-energy cosmic rays. Several refinements and improvements resulted in considerable changes in the predictions as compared to previous studies. In particular, the patterns in the distribution of arrival directions of cosmic rays from astrophysical sources are more pronounced than expected before. The superheavy dark matter scenarios consistent with current data predict less pronounced anisotropy. We developed optimal observables for distinction (exclusion) of these two scenarios. The actual required observational time depends strongly on currently uncertain duty cycle, fluorescent yield and the actual spectrum. In any case, some of the scenarios will be tested even with the limited statistics of the first space-based UHECR detector, TUS, before (or at the time of) the planned launch of JEM/EUSO. However, only JEM-EUSO and/or S-EUSO will be able to detect the SHDM component predicted by the ``minimal'' fit of the AGASA data consistent at 95\% CL with the spectrum and with photon limits. To observe, at the 95\% CL, the patterns correlated with cosmic large-scale structure, one needs $\sim 30$ events in the full-sky sample with energies $E>7 \cdot 10^{19}$~eV. If these patterns show up at high confidence with lower statistics, this might indicate either significant underestimation of the particle energies or a problem in theoretical understanding of the origin and/or propagation of ultra-high-energy cosmic particles. One may try to use these patterns as a rough but independent tool for the energy calibration, quite important for space-based detectors. \ack The authors are indebted to D.~Gorbunov, G.~Rubtsov and D.~Semikoz for discussions. We acknowledge the use of online tools~\cite{XSC,LEDA,GalacticExtinction}. This work was supported in part by the grants RFBR 07-02-00820 (OK and ST) and NS-7293.2006.2 (government contract 02.445.11.7370, OK and ST) and by the Russian Science Support Foundation fellowship (ST). Simulations of the cosmic-ray propagation have been performed at the computer cluster of the Theoretical Division of INR RAS. \appendix | 7 | 10 | 0710.1382 |
0710 | 0710.0384_arXiv.txt | Galactic winds from starbursts and Active Galactic Nuclei (AGN) are thought to play an important role in driving galaxies along the starburst-AGN sequence. Here, we assess the impact of these winds on the CO emission from galaxy mergers, and, in particular, search for signatures of starburst and AGN-feedback driven winds in the simulated CO morphologies and emission line profiles. We do so by combining a 3D non-LTE molecular line radiative transfer code with smoothed particle hydrodynamics (SPH) simulations of galaxy mergers that include prescriptions for star formation, black hole growth, a multiphase interstellar medium (ISM), and the winds associated with star formation and black hole growth. Our main results are: (1) Galactic winds can drive outflows of masses $\sim$10$^8$-10$^9$\msun which may be imaged via CO emission line mapping. (2) AGN feedback-driven winds are able to drive detectable CO outflows for longer periods of time than starburst-driven winds owing to the greater amount of energy imparted to the ISM by AGN feedback compared to star formation. (3) Galactic winds can control the spatial extent of the CO emission in post-merger galaxies, and may serve as a physical motivation for the sub-kiloparsec scale CO emission radii observed in local advanced mergers. (4) Secondary emission peaks at velocities greater than the circular velocity are seen in the CO emission lines in all models, regardless of the associated wind model. In models with winds, however, these high velocity peaks are seen to preferentially correspond to outflowing gas entrained in winds, which is not the case in the model without winds. The high velocity peaks seen in models without winds are typically confined to velocity offsets (from the systemic) $\la$1.7 times the circular velocity, whereas the models with AGN feedback-driven winds can drive high velocity peaks to $\sim$2.5 times the circular velocity. | Observed relationships in galaxies between central black hole mass and stellar mass (e.g. Magorrian et al. 1998), velocity dispersion (i.e. the $M_BH$-$\sigma$ relation; e.g. Gebhardt et al. 2000; Ferrarese \& Merritt 2000), or galaxy structural properties (e.g. the black hole fundamental plane: Hopkins et al. 2007a,b) indicate a co-eval nature in supermassive black hole growth and star formation in galaxies. These results have prompted a number of investigations in recent years to quantify this apparent self-regulation in star formation and black hole growth in galaxies, and their relationship to the formation and evolutionary history of the host system (e.g. Kauffmann \& Haehnelt, 2000; Hopkins et al. 2006a,b; 2007c,d). Over the last two decades, observations of local galaxies have painted a compelling picture in which galaxy mergers provide the link between massive starbursts and central black hole growth and activity. ULIRGs (Ultraluminous Infrared Galaxies), for example, are a class of starburst galaxies with elevated infrared luminosities (\lir $\geq$ 10$^{12}$\lsun) which typically show signs of interactions (e.g. Downes \& Solomon, 1998; Sanders et al. 1988a; Scoville et al. 2000). While the intense infrared luminosity in these sources certainly owes in large part to the merger induced starbursts (e.g. Sanders et al. 1988a,b; Sanders \& Mirabel, 1996), in some cases a contribution from a buried active galactic nucleus (AGN) may be non-negligible. For example, many ULIRGs show spectral energy distributions (SEDs), infrared color ratios (e.g. F[25/60$\mu$m]), polycyclic aromatic hydrocarbon emission (PAH) deficits, and emission line fluxes (e.g. [Ne V] at 14.3$\mu$m) consistent with central AGN activity (e.g. Armus et al. 2004, 2006; Farrah et al. 2003; de Grijp et al. 1985). Indeed, 35-50\% of ULIRGs with luminosity above \lir of 10$^{12.3} L_{\sun}$ show optical and NIR spectra consistent with AGN activity (Kim, Veilleux \& Sanders, 2002; Tran et al. 2001; Veilleux, Kim \& Sanders, 1998). These results suggest that these merging systems are simultaneously undergoing a massive star formation and central black hole growth phase. Observed similarities such as these between starburst galaxies, ULIRGs and quasars prompted Soifer et al. (1987) and Sanders et al. (1988a,b) to propose an empirically derived evolutionary sequence which connects galactic starbursts to quasars through galaxy mergers. In this picture, the fueling of the central black hole and AGN growth phase is realized during the major merger when gaseous inflows trigger nuclear starbursts as well as central black hole growth. The dusty galaxy transitions from a cold, starburst dominated ULIRG, to a warm, AGN-dominated ULIRG (where 'cold' and 'warm' refer to F[25/60$\mu$m] ratios), and, as supernovae and stellar winds clear the obscuring gas and dust, to an optical quasar. In this sense, galaxy mergers provide a unique laboratory for studying the possible co-evolution of starbursts, black hole growth and activity, and spheroid formation. Recent models have provided a theoretical foundation and further evidence for a merger-driven starburst-AGN connection in galaxies. Specifically, simulations by Springel, Di Matteo \& Hernquist (2005a) have shown that galaxy mergers can fuel large-scale gaseous inflows (e.g. Barnes \& Hernquist, 1991, 1996) which trigger nuclear starbursts (Mihos \& Hernquist, 1996; Springel et al. 2005a) as well as promote central black hole growth (Di Matteo, Springel \& Hernquist, 2005). Subsequent winds associated with the growth of central black holes can lift the veil of obscuring gas and dust, and along several sightlines produce a quasar with comparable lifetimes, luminosity functions, and observed $B$-band and X-ray properties to those observed (Cox et al., 2006b; Hopkins et al. 2005a-d; 2006a-d). The merger remnants quickly redden owing to gas depletion and the impact of feedback from star formation and black hole growth (e.g. Springel et al. 2005b) and resemble elliptical galaxies in their kinematic and structural properties. Indeed, the population of stars formed during the starbursts accounts for the central ``excess light'' seen in ongoing mergers (e.g. Rothberg \& Joseph 2004, 2006) and provides a detailed explanation for the luminosity profiles of old ellipticals (e.g. Kormendy et al. 2007; see Mihos \& Hernquist 1994a; Hopkins et al. 2007e,f,g). A consensus picture has thus been borne out from these observations and simulations over the last two decades in which so called 'feedback' processes associated with winds from star formation and black holes act to self-regulate the growth of both the stellar and black hole masses in galaxies (e.g. Fabian, 1999; Silk \& Rees, 1998). Indeed, the effects of starburst and AGN feedback-driven winds have been observed in local galaxies (e.g. Heckman et al. 2000; Martin, 2005; Rupke, Veilleux \& Sanders, 2005a-c; Rupke \& Veilleux, 2005; Tremonti, Moustakas \& Diamond-Stanic, 2007), as well as those at high-\z \ (e.g. Narayanan et al. 2004; Pettini et al. 2002; Shapley et al. 2003) by way of absorption line outflows. However, the direct effect of feedback processes (especially from the highly efficient central AGN) on the emission properties of galaxies is not yet well characterized (see Veilleux, Cecil \& Bland-Hawthorn, 2005 for an extensive review and associated references). In this sense, it is important to relate theories of galaxy formation and evolution which incorporate physically motivated models of feedback processes to observational signatures of winds across the electromagnetic spectrum. In particular, observations of molecular gas in galaxies have proven valuable in characterizing the physics related to nuclear star formation and central AGN fueling in galaxy mergers as the molecular gas serves as fuel both for star formation as well as the central black hole(s). For example, interferometric observations of molecular gas in ULIRGs show that most typically harbor $\sim$10$^{10}$\msun of molecular gas within the central 1.5 kpc (Scoville et al. 1986; Bryant \& Scoville, 1999). Moreover, high resolution maps of dense molecular gas at submillimeter wavelengths in ULIRGs have revealed kinematic structures of double nuclei in mergers (Scoville, Yun \& Bryant, 1997), bar-driven inflows (Sakamoto et al. 2004), and the density structure of gas fueling the central AGN (Iono et al. 2004). With the increased spatial resolution and sensitivity afforded by the latest generation of (sub)mm-wave interferometers (e.g. the SMA, CARMA, PdBI), a number of recent observations have been able to pioneer investigations as to the effects of galactic winds on molecular gas emission in galaxies (e.g. Iono et al. 2007; Sakamoto, Ho \& Peck, 2006; Walter, Wei\ss \ \& Scoville, 2002). Detections such as these are expected to become more routine in upcoming years as the ALMA interferometer becomes available. An important complement to these current and forthcoming observations of molecular line emission from starburst galaxies and AGN are physical models which directly relate CO emission properties to galactic scale winds. In this context, it is our aim to investigate the role that galactic winds can play on CO emission properties from starburst galaxies and AGN via numerical simulations. In particular, we focus on specific signatures of winds imprinted on CO morphologies and emission line profiles. In this work, we present self-consistent radiative transfer calculations for the emission properties of CO molecular gas in gas-rich galaxy mergers which account for the winds associated with both star formation and black hole growth. In \S~\ref{section:numericalmethods}, we describe the hydrodynamic simulations, and radiative transfer methodology. In \S~\ref{section:spiral} we provide an example of our methods by applying our radiative transfer calculations to a model of a star-forming disk galaxy. In \S~\ref{section:morphology}, we discuss the effect of winds on observed CO morphologies. We explore the response of emission line profiles to winds in \S~\ref{section:lineprofiles}, present a broader discussion of these results with respect to observations in \S~\ref{section:observations}, and summarize in \S~\ref{section:conclusions}. Throughout this paper, we assume a $\Lambda$CDM cosmology with $h$=0.7, $\Omega_\Lambda$=0.7, $\Omega_{\rm M}$=0.3. | \label{section:conclusions} We have combined 3D non-LTE radiative transfer calculations with SPH simulations of galaxy mergers in order to investigate the effects of galactic-scale winds on the molecular line emission in starburst galaxies and AGN. We find that galactic winds are a natural result of merger-induced star formation and black hole growth. These winds can entrain molecular gas of order $\sim$10$^8$-10$^9$ \msun which imprints generic signatures in both the CO morphology as well as unresolved emission line profiles. The specifics of the morphological and emission line indicators of molecular outflows depend on physical parameter choices within the galaxy merger models. In particular, the energy source (i.e. BH accretion or star formation), as well as the merger orientation can vary the strength, direction, and duration of molecular outflows in emission contour maps, as well as the velocity separation of high velocity peaks in multi-component emission lines. Many of the results including the halflight radius as a function of excitation, or the velocity offsets of high velocity peaks in the emission lines vary enough between physical parameter choices such that these models may be used to constrain physical origins for observed CO morphologies and line profiles. In detail, we find the following: \begin{enumerate} \item Molecular outflows entrained in AGN feedback and starburst-driven winds can be directly imageable via CO emission line mapping. These types of outflows have been recorded in at least some local systems (e.g. Iono et al. 2007). These outflows will be detectable at cosmological distances given the predicted spatial resolution of ALMA. \item Molecular outflows entrained in winds driven by AGN feedback are typically longer lived than those entrained in starburst driven winds. This owes to the relative strength of AGN feedback-driven winds versus starburst driven winds. \item Winds from AGN feedback in coplanar merger models are typically more powerful than mergers which occur at more random orientations. This results in outflows in these models being visible for the majority of the $\sim$200 Myr 'active period' studied here. \item The spatial extent of CO emission can be controlled by the presence of galactic winds. The emission is seen to be stratified with transition in all models such that the CO (J=1-0) halflight radius is larger than the radius from higher lying transitions (e.g. CO J=7-6), though the degree of stratification depends on the inclusion of winds. The relatively compact nature of observed CO emission in local mergers (R$_{1/2}$ typically confined to the central kpc; Scoville \& Bryant, 1999) may be a consequence of AGN feedback or starburst-driven winds. \item In all models, high velocity peaks (peaks at velocities greater than the circular velocity) can exist superposed on the post-merger galaxy's broad CO emission line. In models without winds, these peaks owe to random kinematics of molecular gas. In models with winds, these peaks are seen to originate primarily from gas entrained in outflow, at least during the period of peak black hole accretion/star formation. \item For the models studied here (Table~\ref{table:ICs}), high velocity peaks driven by random kinematics do not typically appear at velocity offsets (from systemic) greater than $\sim$1.7 times the circular velocity of the post-merger galaxy. In contrast, peaks entrained in AGN feedback-powered winds can be driven to velocities near 2.5 times the circular velocity. The centroid velocities of the simulated lines are typically (1$\sigma$) within $\sim$20 \kms of the true systemic velocity, and can generally be used as a reliable substitute for the systemic velocity in these models. Thus the above results hold true when measuring the velocity offsets of high velocity peaks with respect to line centroids as well. \end{enumerate} | 7 | 10 | 0710.0384 |
0710 | 0710.2518_arXiv.txt | { {} {Using the high-level Balmer lines and continuum, we trace the density structure of two magnetospheric disk segments of the prototypical Bp star $\sigma$\,Orionis\,E (B2p) as these segments occult portions of the star during the rotational cycle. } {High-resolution spectra of the Balmer lines $\ge$H9 and Balmer edge were obtained on seven nights in January-February 2007 at an average sampling of 0.01 cycles. We measured equivalent width variations due to the star occultations by two disk segments 0.4 cycles apart and constructed differential spectra of the migrations of the corresponding absorptions across the Balmer line profiles. We first estimated the rotational and magnetic obliquity angles. We then simulated the observed Balmer jump variation using the model atmosphere codes \mbox{\it synspec/circus} and evaluated the disk geometry and gas thermodynamics. } {We find that the two occultations are caused by two disk segments. The first of these transits quickly, indicating that the segment resides in a range of distances, perhaps 2.5$-$6R$_{*}$, from the star. The second consists of a more slowly moving segment situated closer to the surface and causing two semi-resolved absorbing maxima. During its transit this segment brushes across the star's ``lower" limb. Judging from the line visibility up to H23-H24 during the occultations, both disk segments have mean densities near 10$^{12}$ cm$^{-3}$ and are opaque in the lines and continuum. They have semiheights less than ${\frac 12}$R$_{*}$, and their temperatures are near 10\,500\,K and 12\,000\,K, respectively. } {In all, the disks of Bp stars have a much more complicated geometry than has been anticipated, as evidenced by their (sometimes) non-coplanarity, de-centerness, and from star to star, differences in disk height. } | Over the last several years the interest in the class of magnetic Bp stars has spurred the accumulation of datasets at many wavelengths over the stars' rotational/magnetic cycles. These datasets support the view that the circumstellar environment is the site of a balance of physical processes that channels the wind toward the magnetic plane. If the magnetic field is strong, that is if the circumstellar magnetic energy density exceeds the wind energy density, the star's wind is channeled into the magnetic plane where a shock is formed and the cooled shock settles to form a stable disk. The magnetic axis is typically inclined with respect to the rotational axis. Then, as different portions of the disk wobble in front of the star during the rotational cycle, they produce alternate absorption and emission components in the H$\alpha$, UV resonance lines, and low-excitation metallic lines (e.g., Bohlender et al. 1987, Shore 1987, Bolton 1994). To complicate the picture, it appears that slow infall of disk matter is responsible for producing inhomogeneous metal-poor patches or a belt along the magnetic equator (e.g., Khokhlova et al. 2000, Groote 2003). Also, although there remains a theoretical difficulty with the fractionation hypothesis for helium (Krti\v{c}ka et al. 2006), enough neutral helium atoms in fact seem likely to separate from the polar wind, return to the surface, and accumulate there as helium caps. Recent theoretical refinements in the picture (e.g., Preuss et al. 2004, Townsend \& Owocki 2005; hereafter TO05) suggest that the wind will settle onto low equipotential surfaces determined by the balance of radiative, gravitational, magnetic, and centrifugal forces. These surfaces reside primarily, though not exclusively, in the disk plane. Because the disk is likely to have a nonaxisymmetric density distribution, we will refer to the accumulations near the plane responsible for time-variable spectral features as disk segments. TO05 have presented a ``rigidly rotating magnetospheric" {\it ab initio} model to explain physical properties of these segments for the case of $\sigma$\,Ori\,E and it can be applied to the cases of other Bp stars with arbitrary magnetic inclinations and viewed from any rotational inclinations. TO05 find that the disk plane is not precisely perpendicular to the magnetic polar axis for intermediate magnetic obliquities, including those they obtained for $\sigma$\,Ori\,E. The spectroscopic anomalies of $\sigma$\,Ori\,E (HD\,37479, B2\,Vp; Lesh 1968) were first noticed by Berger (1956) and discussed further by Osmer and Peterson (1972). This star is the prototype and most studied of the helium-strong subclass of Bp stars. Walborn (1974) discovered H$\alpha$ emissions in this star's spectrum that varied on a timescale of hours. This variability was subsequently established to be associated with the rotation of a frozen-in disk, as described above. Hesser et al. (1977) determined a period of 1.19081 days based on several nights of photometric observations distributed over a few years. Subsequently, various studies (Landstreet \& Borra 1978, Bohlender et al. 1987) determined that the magnetic variations (suggestive of an approximately dipolar field with polar field strength B$_p \approx$ 10\,kG) are consistent with the photometric period. Various periods over the range 1.19081$-$1.19084 days have since been determined by several authors from analyses of variations of photospheric and circumstellar H$\alpha$ line emissions (Hesser, Moreno, \& Ugarte 1977, Reiners et al. 2000, hereafter R00; Townsend, Owocki, \& Groote 2005; hereafter TOG05). Groote \& Hunger (1976, 1977 hereafter GH76 and GH77) discussed the variations of high-level Balmer lines based on only photographic spectra. The GH76 data indicated that two disk occultations of the star, about 0.4 cycles apart, cause high-level Balmer line absorptions visible at these times. At the same times, the intervening disk material absorbs flux just shortward of the Balmer edge. Although the analysis of line absorptions do not generally allow us to infer the extent and total volume of such a corotating structure, these absorptions do yield considerable information about its three dimensional geometrical character. Smith \& Groote (2001, hereafter SG01) combined strength and shape information of many ultraviolet metallic lines obtained by the {\it International Ultraviolet Explorer} around the rotation cycle to exploit this potential. Their study provides some of the first estimates of the physical parameters of the disk gas, including disk temperature, column density, and areal coverage of the star. The limitations of this study were the low signal-to-noise ratio and the paucity of phase sampling of the archival database. What has been needed is a study that includes complete rotational phase coverage of absorption lines sensitive to electron density. To respond to this need, we have obtained complete phase coverage of the hydrogen lines from H9 (H$\zeta$) to the Balmer edge over the entire rotation period of $\sigma$\,Ori\,E. Such datasets have the potential to better elucidate areal coverage and density structure of the absorbing disk segments than has previously been possible. | We have used high-level Balmer absorption data finely sampled around the rotation cycle to show that the distribution of plasma in the magnetosphere of $\sigma$\,Ori\,E is more complicated than previously believed. Most especially, the secondary occultation is comprised of two semi-resolved events. In addition, there seem to be at least two brief appearances of absorbing matter well beyond the disk plane. The disk plane is dominated by two segments which occult portions of the star some 0.4 cycles apart. We have no direct information from absorption profiles alone on the distribution of matter at other azimuths in the plane and so look forward to future analyses based on the H$\alpha$ and/or H$\beta$ emissions over the cycle. The occulting disk segments are sections of a ring which occult the star as the frozen disk corotates around the star. The segment causing the primary occultation lies further from the star (between limits of $\approx$2.5R$_{*}$ and 6R$_{*}$) than the opposing segment. There are at least three arguments that point to the disk segments lying at different orbital radii. First, gas in the segment causing the primary occultation has a lower temperatures (10\,500\,K, compared to 12\,000\,K). This can be inferred first from the lack of an enhanced He\,I feature and also from the absence of weak absorptions of metallic-line blends referred to in $\S$4.3. Second, the acceleration of spectral migrating components across the line profile is higher for this segment. Third, our occultation analysis program indicates that one cannot reconcile the radii of the two segments even if one attributes the broad secondary absorption maximum as being due to viewing two occultations of a single warped disk segment. Perpendicular to the magnetic disk, the semiheights of the two segments are some 0.3$-$0.45R$_{*}$. It is important to stress that the disk of $\sigma$\,Ori\,E has a smaller semiheight than disks for other disk-harboring Bp stars that we have analyzed with similar modeling techniques (e.g., SG01, Smith et al. 2006). It is tempting to speculate that a disk should be more compressed toward the plane due to the pressure provided by a large magnetic field (e.g., Babel \& Montmerle 1997). However, this cannot be the full story because the most relevant quantity, the ratio ``$\eta$" of magnetic to wind energy densities is actually larger for 36\,Lyn than for $\sigma$\,Ori\,E and yet the disk of the former has a somewhat larger height (Smith et al. 2006). Might other considerations, such as the larger radial extent of the disk (e.g. in the case of 36\,Lyn), have a bearing on this question? In addition to these dimensional properties, we find that a phase discrepancy of +0.018 cycles exists between the phase at which the central absorption transits the center of the hydrogen line and the phase of EW maximum. This agrees with a similar phase difference found by GH77 and calls into definition of the star's period when determined from different observation modes. The secondary occultation is distinguished by the passage of two weak lobes, the phases of which are consistent with the semi-resolved structure of the light curves in the literature. The interpretation of this structure is not clear, but it may be related to a disk ``warping" or a corrugation of the disk plane predicted in the TO05 models. When we view the disk edge-on through the magnetic plane, the Balmer line absorptions are opaque, and thus the absorptions of the high-level members of this series increase. For H9 the optical depth is about 100, while it is $\approx 10$ for H20, and $\approx 2$ at the Balmer edge. As the disk ingresses and egresses there is an absence of an absorption tail in the equivalent width curves. This confirms the theoretical expectation (e.g., TO05) that the density scale height is very small compared to the stellar radius. This implies in turn that our homogeneous slab models of the disk and the total mass estimate of about 1$\times$10$^{-10}$ M$_{\sun}$ are only rough approximations. How would hypothetical data of indefinitely high quality (signal-to-noise, spectral resolution, and cadence) improve on our analysis of the disk of $\sigma$\,Ori E? We can think of the following applications of enhanced quality data: One improvement would be in the mapping of vertical density distribution and indeed the separation of the two density maxima associated with the secondary occultation. This would result from tracing details of the absorption subfeatures as they migrate across the line profiles. An improvement in our understanding of the vertical stratification should indirectly improve out understanding of whether the disk is warped and the inner edge of the secondary occultation-causing segment extends almost to the star or not. It will also clarify the cause of the phase delays that both we and GH76 have found among various hydrogen lines and with respect to the continuum. A resolution of this issue would allow an improved determination of ephemerides from flux curves derived from different diagnostics, thereby leading to a definitive determination of the period and perhaps its first derivative. Second, the changes in the highest-level lines detected during the ingress and egress of the occultations can be used in principle to further map the vertical density distribution of the disk segments because of the different optical depth regimes in which the various lines are formed. Third, highly accurate difference profiles should allow one to search for asymmetric absorption features in the EW curves that could indicate the rate at which disk particles leak from the disk either through the inner or outer edge. Fourth, increased data quality improvements (assuming a stable spectrograph and atmospheric transmission) can lead to the observation of subtle changes in disk structure. We suggest that these might occur on timescales of one to a few months or longer. Finding such changes could allow one ultimately to correlate them with the expected ``break out" of disk matter through the outer edge (ud-Doula, Townsend, \& Owocki 2006), possibly associated with X-ray flares from this star, (Groote \& Schmitt 2004, Sanz-Forcada, Franciosini, \& Pallavicini 2004). | 7 | 10 | 0710.2518 |
0710 | 0710.2668_arXiv.txt | We report the detection of 10 new X-ray filaments using the data from the {\sl Chandra} X-ray satellite for the inner $6^{\prime}$ ($\sim 15$ parsec) around the Galactic center (GC). All these X-ray filaments are characterized by non-thermal energy spectra, and most of them have point-like features at their heads that point inward. Fitted with the simple absorbed power-law model, the measured X-ray flux from an individual filament in the $2-10$ keV band is $\sim 2.8\times10^{-14}$ to $10^{-13}$ ergs cm$^{-2}$ s$^{-1}$ and the absorption-corrected X-ray luminosity is $\sim 10^{32}-10^{33}$ ergs s$^{-1}$ at a presumed distance of 8 kpc to the GC. We speculate the origin(s) of these filaments by morphologies and by comparing their X-ray images with the corresponding radio and infrared images. On the basis of combined information available, we suspect that these X-ray filaments might be pulsar wind nebulae (PWNe) associated with pulsars of age $10^3 \sim 3\times 10^5$ yr. The fact that most of the filament tails point outward may further suggest a high velocity wind blowing away form the GC. | The Galactic center (GC) is the only place where we can observe parsec details of various interaction in and around the Galactic nucleus. Advances in this research frontier rely primarily on observations at radio, infrared and X-ray wavelengths, because the optical band suffers seriously from a considerable extinction with an $Av\sim 30$ (e.g., Becklin, Matthews, Neugebauer \& Willner 1978). One of the most important discoveries in the GC region is perhaps the presence of many structured non-thermal radio filaments (NTFs) (e.g., Yusef-Zadeh, Morris, \& Chance 1984; Morris \& Serabyn 1996; LaRosa, Kassim, Lazio, \& Hyman 2000). While these non-thermal radio filaments have been intensively studied, their origins and implications on the underlying physical processes around the GC region remain largely unclear. In the current analysis of X-ray filaments around the GC region, these X-ray emitting particles are usually expected to be fairly close to their acceleration zone and evolve very rapidly in time. Thus, the X-ray study of the same region would be essential to probe the origin of these energetic particles. The {\sl Chandra} Galactic center Survey (CGS), with its unprecedented high spatial resolution of $\sim 0.5''$ and moderate spectroscopy capability, has already revealed remarkable X-ray structures (including thousands of X-ray bright point sources and some filaments, as well as clumps of diffuse emission) within the central $200$ pc of our Galaxy (e.g., Wang, Lu, \& Lang 2002a). In this paper, we mainly concentrate on the nature of those thread-like linear structures or filaments as observed in X-ray bands. Up to this point within $15^\prime$ ($\sim 37$ pc at 8 kpc) from \hbox{Sgr~A$^*$} where a $\sim 4\times 10^6M_{\odot}$ black hole resides inside a compact region of less than $\sim 1$AU (e.g., Shen et al. 2005), 5 X-ray filaments have been studied in details (see also Table 1). For example, G359.95-0.04, a comet-like filament at $\sim 0.32$ pc north of \hbox{Sgr~A$^*$}, is thought to be a ram-pressure confined pulsar wind nebula (PWN) (Wang et al. 2006). Another prominent filament, G359.89-0.08 (SgrA-E), at $\sim 7$ pc southeast of \hbox{Sgr~A$^*$}, was first noticed by Sakano et al. (2003) and an interpretation of a possible PWN origin was discussed in details by Lu, Wang \& Lang (2003). An alternative picture of G359.89-0.08 as a source of synchrotron emission from relativistic particles accelerated by a shock wave of W28 SNR was suggested recently by Yusef-Zadeh et al. (2005); in that same paper, they also detected a new X-ray filament G359.90-0.06, which coincides spatially with a radio filament at $\sim 5.8$ pc southwest of \hbox{Sgr~A$^*$}. They explored the mechanism of an inverse Compton scattering (ICS) for the X-ray emission of G359.90-0.06. Another 2 filaments were found in more extensive regions. G0.13-0.11 in the Radio Arc region was first reported by Wang, Lu \& Lang (2003) and was suspected to be also a PWN. Of particular interest is the X-ray filament G359.54+0.18; it associates with only one strand of the obviously bifurcated radio threads (e.g., Yusef-Zadeh et al. 2005; Lu et al. 2003). A common feature shared in the X-ray energy spectra of these filaments is that they all appear to be non-thermal. It should be noted however that any thermal component to these sources would likely be completely absorbed and unobservable, given the high foreground column density. In these earlier investigations, pulsar wind nebulae (PWNs) and supernova remnants (SNRs) seem to offer natural explanations for the appearance of such X-ray filaments. Indeed, it is believed that a considerable number of supernovae should have happened in the GC region (e.g., Figer et al. 1999, 2004; Wang et al. 2006). One would naturally expect to find some of their end-products such as pulsars and SNRs in the GC region. However, no radio pulsars have yet been found within $\sim 1^\circ$ of the GC (Wang et al. 2006). This might be caused by difficulties in radio observations (Cordes \& Lazio 1997; Johnston et al. 2006). Seeking observational clues in X-ray bands might shed new light to the search for radio pulsars embedded in the GC region. Another tempting idea is to use these X-ray filaments as potential tracers for the magnetic field and gas dynamics around the GC region, since the magnetic fields should have played a significant role in producing such non-thermal spectra and the thread-like shapes of these filaments and the gas motion is usually coupled with the magnetic fields. (e.g., Chevalier 1992; Boldyrev \& Yusef-Zadeh 2006). Magnetic fields exist on all scales of the Galaxy as well as in other spiral galaxies and generally trace out spiral patterns on large scales (e.g., Beck \& Hoernes 1996; Fan \& Lou 1996; Zweibel \& Heiles 1997; Wielebinski 2005; Ferriere 2001, 2007), and great progress has been made in measuring them and inferring their influence by various means. For example, the observed high-energy cosmic ray anisotropy at a few times $10^{-3}$ (e.g., Amenomori et al. 2006) might be physically related to large-scale structures of Galactic magnetic fields and inhomogeneous cosmic ray source distribution. Using diffuse synchrotron radio emissions at 74 and 330 MHz frequencies produced by relativistic cosmic-ray electrons and the magnetic field around the Galactic center, LaRosa et al. (2005) inferred a weak magnetic field of order of $\sim 10 \mu$G on size scales $\gsim 125''$ based on a minimum-energy analysis. This is about 2 orders of magnitude lower than $\sim 1$ mG usually estimated for the GC region. Very recently, Cowin \& Morris (2007) argued that the assumption of LaRosa et al. (2005) that the magnetic field and cosmic rays are in a minimum-energy state across this region is unlikely to be valid. According to their model estimates, the mean magnetic field is at least 100 microgauss on a scale of several hundred parsecs and peaks at approximately 500 microgauss at the center of the diffuse nonthermal source (DNS). This is an important issue to be settled for the GC magnetic environment. Most of the GC radio NTFs are found to be perpendicular to the Galactic plane, implying a local poloidal magnetic fields of about milli-gauss strengths (e.g., Yusef-Zadeh \& Morris 1987). However, recent observations revealed that the GC magnetic field may be more complex than a simple globally ordered dipolar field (e.g., LaRosa et al. 2004; Nord et al. 2004). It might be possible for X-ray NTFs to also provide clues of the configuration of the local magnetic field as well as the interaction between it and the ambient gas flow. We report morphological and spectral properties of another 10 newly discovered X-ray filaments within a region of $6^\prime .5 \times 6^\prime .5$ surrounding the \hbox{Sgr~A$^*$} (roughly corresponding to a projected sky area of $\sim 15$ pc by $\sim 15$ pc at a presumed distance of 8 kpc to the GC). Their plausible physical origins are discussed in section 4. All error bars in the X-ray spectrum parameter measurements are at the 90\% confidence level, and we express the fitted parameters in the format of $y\ (y_l,\ y_u)$, where $y$ is the best fit value while $y_l$ and $y_u$ are the lower and upper limits of the 90\% confidence interval, respectively. For all the images in the paper, north is up and east is to the left. We shall adopt a distance of 8 kpc from the solar system to the GC throughout this paper. We note that during the review process of this manuscript, Muno et al. (2007) submitted a paper also on X-ray filaments around the Galactic center. | \subsection{\label{subsec:origins}Properties and possible origins of GC X-ray filaments } While X-ray photon numbers may not be high enough to constrain the exact shapes of the energy spectra, all spectra appear featureless except for filament F3 showing weak iron line features (see Section 3.3) and can be well fitted with power-law models. Fitting of some of these energy spectra with thermal emission models is acceptable statistically; nevertheless, this always gives quite high temperatures, i.e., $>10$ keV. We therefore incline to the view that X-ray emissions from these filaments are non-thermal in nature. Table 2 sums up the inferred parameters for these 15 non-thermal X-ray filaments in the inner $15^{\prime}$ around the GC. In addition to the 10 filaments studied in this paper, we also include the other 5 filaments, namely G359.89-0.08, G359.90-0.06, G359.95-0.04, G359.983-0.046, and G0.13-0.11, reported and studied earlier in the literature. Most of their hydrogen column density $N_H$ are of the order of $\sim 10^{22}-10^{23}~{\rm cm}^{-2}$, consistent with other $N_H$ estimates around the GC region. Most of their photon indices $\Gamma$ fall within the range of $\sim 1-2.5$. The chance of these X-ray filaments being background extragalactic sources is very small according to the spectral and morphological properties. For instance, it would be very difficult to explain the linear filamentary morphology using the hypothesis of extragalactic origins. Therefore, these X-ray filaments are most likely unique objects around the GC region. As already discussed in Section 1, there were suggestions that these X-ray filaments may be ram-pressure confined PWNe (Wang et al. 2003, 2006), or synchrotron emissions from MHD shocks associated SNRs or emissions resulting from inverse Compton scattering (Yusef-Zadeh et al. 2005; Figer et al. 1999). The non-thermal X-ray emission mechanisms may be either synchrotron emission or inverse-Compton scattering. Magnetohydrodynamic (MHD) relativistic pulsar winds (Michel 1969; Goldreich \& Julian 1970; Kennel \& Coroniti 1984a, b; Lou 1996, 1998) and MHD shock interactions of magnetized outflows (e.g., Yu \& Lou 2005; Yu et al. 2006; Lou \& Wang 2006, 2007) with the interstellar medium (ISM) in SNRs and PWNe could provide high-energy electrons needed in these two radiation mechanisms (e.g., Sakano et al. 1993; Lu et al. 2003; Wang et al. 2006). If these X-ray filaments are SNRs, their elongations would probably represent MHD shock fronts and therefore, one would not expect to see a tendency of spectral softening along a filament. The two bright filaments F3 and F10 both show evidence for such a softening tendency. The fact that most of these filaments have point-like sources at the heads also againsts the SNR origin. On the other hand, as nonthermal X-ray emission is only detected in several SNRs within the entire Galaxy, it would be highly unlikely that there are so many nonthermal SNRs around the GC region. For this reason, we would argue that most of these X-ray filaments are not SNRs. Observed properties of these X-ray filaments may be more consistent with those of PWNe. Typical features of a PWN are: non-thermal X-ray spectrum, with photon index of $1.1-2.4$ and a X-ray luminosity $L_x$ range from $4\times 10^{32}$ to $2\times 10^{37}~{\rm erg}~{\rm s}^{-1}$ in $0.2-10$ keV band (e.g., Gaensler \& Slane 2006; Kaspi, Roberts \& Harding, 2006). The $\Gamma$ and $L_x$ of these 10 filaments are consistent with the values of a PWN. The existence of point-like X-ray sources as indicated by the image study also tends to favor a PWN scenario. We may estimate the ages of the putative pulsars with the X-ray luminosities of these X-ray filaments. Li et al. (2007) studied statistically the nonthermal X-ray emission from young rotation powered pulsars and PWNe. They noted that there exists a correlation between the pulsar age $\tau$ and the $2-10$ keV PWN luminosity $L_{x,pwn}$, which can be expressed as $L_{x,pwn}=10^{41.7}\tau^{-2.0\pm 0.3}$. The X-ray luminosities of these 10 filaments are in the range of 0.2-2.2$\times 10^{33}$ erg s$^{-1}$. Using this empirical formula, ages of these putative pulsars are possibly between $\sim 10^3$ to $3\times10^5$ yr. However, given the dispersion about the above empirical relationship (Li et al. 2007), the estimate may be uncertain probably by a factor of 10. Since the ages of pulsars with bright PWNe are usually younger than a few tens of thousand years, one may doubt if a pulsar at the age of several $10^5$ years can produce a detectable X-ray nebula. However, the PWN of a relatively old pulsar can be enhanced in surface brightness and thus become detectable if the pulsar wind materials are confined to one direction. PSR B0355+54 is $\sim 5.6\times10^{5}$ yr old. It converts $\sim 1$\% of its spin-down luminosity to the cometary-like X-ray nebula (e.g., Tepedelenlio\v{g}lu \& \"{O}gelman 2007). The old pulsar PSR B1929+10 ($\tau\sim3\times10^6$ yr) also converts $2.1\times10^{-4}$ of its spin-down luminosity $\sim 3.9\times10^{33}$ erg s$^{-1}$ into the emission of the cometary nebula (e.g., Becker et al. 2006). The X-ray filaments identified in the GC region are similar to these two systems and thus probably powered by pulsars. Now we discuss whether the number of X-ray filaments, if identified with PWNe, would be consistent with the estimated star formation rate in the GC region. According to Figer et al. (2004), the star formation rate at the GC is about $10^{-7} M_{\odot}\hbox{ yr}^{-1}\hbox{ pc}^{-3}$ which is some 250 times higher that the mean star formation rate in the Galaxy. In the field of view of our Fig. 1, we take a radius of about $7$ pc and estimate the star formation rate to be $1.4\times 10^{-3}M_{\odot}\hbox{ yr}^{-1}$. If the mean mass of a star is 10$M_{\odot}$, the frequency of supernova explosions would be $1.4\times 10^{-4}$ per year, leading to about 40 pulsars in the field of Fig. 1 younger than $\sim 3\times 10^5$ yr as estimated above. This number is roughly consistent with the 15 candidate PWNe identified in the field. \subsection{\label{subsec: F10_radio}G0.029-0.06 (F10) and its radio counterpart} Filament F10 bears certain unique features to be noted here. First, it has the longest linear structure with the entire image slightly bent towards the northeast, more or less like an arc. Second, it is the farthest away from \hbox{Sgr~A$^*$} and thus has much less contamination from the strong diffuse X-ray emission of Sgr A. Third, there is an obvious 20cm radio NTF coincident spatially with X-ray filament F10. The spectral indices for different regions along filament F10 show evidence of spectral steepening from the ``head" to ``tail" (see Table 1). When $N_H$ is fixed at the best fit value $\sim 7\times 10^{23}~{\rm cm}^{-2}$, the $\Gamma$ values for the ``head", ``middle", and ``tail" regions are 1.1(1.0, 1.3), 1.5(1.4, 1.6), and 1.8(1.7, 1.9), respectively. This might suggest an energetic particle flow direction from the southeast (head) to the northwest (tail). A pulsar moving through the magnetized interstellar medium seems to give a plausible explanation of this scenario. Indeed, the morphology of F10 does imply a point source in the ``head" region. The corresponding point spread function (PSF) at G0.029-0.06 is an ellipse with a size of $\sim 2\arcsec\times 4\arcsec$. For an updated X-ray versus spin-down luminosity correlation of rotation powered pulsars, a modified empirical relation is given by equation (3) of Possenti et al. (2002), namely, $\log L_{x,(2-10)}=1.34~\log\dot{E}-15.34$ where $L_{x,(2-10)}$ is the X-ray luminosity in $2-10$keV energy band; using this empirical relation, we would have a $\dot E\sim 10^{36}~{\rm ergs}~{\rm s}^{-1}$. Since PSRs J1747-2958 and B1929+10 convert about 2.5\% and 2.1$\times10^{-4}$ of their spin-down powers to their cometary X-ray nebulae (e.g., Gaensler et al. 2004; Becker et al. 2006), the ratio $L_x$/$\dot{E}$ of F10 ($\sim 10^{-3}$) indicates that the above estimate for $\dot{E}$ is reasonable. The arc-like X-ray morphology of F10 and its coincidence with a radio NTF might be a good indicator of its interaction with the interstellar magnetic field environment of the GC region (Lang et al. 1999; Wang et al. 2002b). Similar to the discussion about G0.13-0.11 by Wang et al. (2002b), we may estimate the magnetic field strength $B$ in the current context. First, the lifetime $\tau$ of synchrotron X-ray emitting particles is given by $\tau\sim (1.3 {\rm~yrs})\epsilon^{-0.5} B_{\rm mG}^{-1.5}$, where $\epsilon$ is the X-ray photon energy in unit of keV (a value of 4 keV is adopted here) and $B_{\rm mG}$ is the magnetic field strength in the filament volume in units of mG. The simulations of Bucciantini et al. (2005) show that the average flow speed in the tail is about 0.8-0.9 $c$. For a sustained X-ray linear structure, we estimate by requiring $\tau \geq L_{obs}/(0.85 c)$. Adopting a characteristic angular length $L_{obs}$ of $47\arcsec$ ($\sim 2 {\rm pc}$), we thus infer a magnetic field strength $B\sim 0.3$mG, similar to those in the bright radio NTFs (e.g, Yusef-Zadeh \& Morris 1987; Lang et al. 1999). We try to outline a few plausible scenarios in the present context and discuss relevant aspects qualitatively. Magnetized neutron stars move with peculiar speeds in the range of a few tens of kilometers per second (a mean space velocities of $\sim 300-400\hbox{ km s}^{-1}$ for young pulsars; Hobbs et al. 2005; Faucher-Gigu\`ere \& Kaspi 2006) to well over one thousand kilometers per second ($\sim 1600\hbox{ km s}^{-1}$) and the surrounding ISM is generally magnetized. Generally speaking, a typical peculiar velocity of a neutron star is supersonic and super-Alfv\'enic in a magnetized ISM. Neutron stars or pulsars have different ranges of surface magnetic field strengths: $10^9-10^{10}$G for millisecond pulsars in binaries, $10^{11}-10^{12}$G for a wide range of pulsars, and $10^{14}-10^{15}$G inferred for several magnetars. Several situations may happen. (1) If a pulsar does not involve an active pulsar wind, its peculiar motion through the surrounding magnetized ISM would sustain an MHD bow shock by its magnetosphere as well as a magnetotail. The faster the pulsar moves, the more linear the system would appear. This is basically like a bullet moving through an air supersonically and generating a Mach cone or wake. Relativistic electrons can be produced at the MHD bow shock and synchrotron emissions can be generated and sustained at the same time. (2) In a binary system, the fast wind (say, with a speed higher than 1000 km s$^{-1}$) from a companion star can blow towards a spinning magnetized pulsar in orbital motion. Here, the situations of a companion fast wind blowing across a magnetized pulsar and a pulsar moving through the ISM with a high speed are more or less equivalent. Again, an MHD bow shock and a magnetotail can form in association with the pulsar system. The stronger the companion wind and the faster the pulsar moves, the more linear the pulsar system would appear. Relativistic electron and/or positrons can be generated and sustained to power synchrotron emissions in the bow shock draped around the pulsar magnetosphere. For such a system, one might be able to detect the presence of the companion by various independent means. (3) For a pulsar emitting an active pulsar wind and with misaligned magnetic and spin axes, spiral forward and reverse shock pairs can be generated in the relativistic pulsar wind as a result of inhomogeneous wind and eventually the pulsar wind is stopped by the ISM through a MHD termination shock (e.g., Lou 1993, 1996, 1998). (4) Case (3) can also happen for a pulsar moving with a high peculiar velocity through a magnetized ISM. (5) Case (3) can also happen for a pulsar in binary orbital motion with the companion blows a powerful wind with a speed higher than 1000 km s$^{-1}$. In both cases of (2) and (5), the center of mass of the binary system may also move with a high speed through the ISM. One can further speculate several possible combinations along this line of reasoning (e.g., Chevalier 2000; Toropina et al. 2001; Romanova, Chulsky \& Lovelace 2005). There are two clumps (referred to as the east and west clumps hereafter) of diffuse X-ray emission surrounding filament F10. Although the west clump is also elongated, it is not called a filament because it contains many substructures. To see if these clumps are physically related to F10, we extract their energy spectra separately (see the middle and bottom panels of Fig. 6). The fitted parameters are listed in Table 3. The much higher absorbing column densities of the two clumps indicate strongly that they are located farther away from us than F10 is, while these two clumps themselves are almost at the same distance (see Table 3). Moreover, their X-ray emissions are very likely powered by the same mechanism, as hinted by the characters of their spectra, which can be fitted well with an absorbed power-law model plus a 6.4 keV emission line. The total emission comes mostly from the photon energy $4-7$ keV band, with a strong 6.4 keV neutral Fe K line. In contrast, the Fe line feature is not present in the spectrum of F10. A possible explanation for this non-thermal, apparently broadened iron line emission at 6.4 keV is the collisions of low-energy cosmic-ray electrons with irons in molecular clouds (e.g., Valinia et al. 2000) or by the radiative illumination from the GC massive black hole that was suggested to be very bright in the past (e.g., Koyama et al. 1989). In conclusion, F10 and the surrounding clumps do not seem to interact directly. \subsection{\label{subsec:B-fields}X-ray NTFs as tracers of the small-scale magnetic fields and gas dynamics} The oritentation of the X-ray filaments provides an opportunity to probe the physics conditions of the GC region. As discussed in section 4.1, the cometary shapes of the X-ray filaments imply that the pulsar wind particles are confined to one direction by the ambient materials. Mechanism shaping the filamentary structure could be ordered magnetic fields (e.g., Yusef-Zadeh \& Morris 1987; Lang, Morris, \& Echevarria 1999) and/or high relative velocity between the pulsars and the surrounding gas (e.g., Wang et al. 1993; Shore \& LaRosa 1999). The magnetic field could be the product of the gas motion, and the magnetic field could also control the motion of the gas (e.g., Heyvaerts et al. 1988; Chevalier 1992). Radio NTFs suggests that the magnetic field is poloidal at large-scale (e.g., Yusef-Zadeh \& Morris 1987; Lang, Morris, \& Echevarria 1999) with some more complex smaller structures around the GC region (e.g., Nord et al. 2004). By looking at the X-ray images shown by Figures 1, 2 and 3, there seems to be a tendency that the PWN tails point away from the GC, indicating that the pulsar wind particles are blown outward. This might imply the presence of a Galactic wind of hot plasma blowing away from the center, given the high star formation rate (and so plenty of hot gas) in this region. Pulsars may have typical peculiar velocities of $\sim 400\hbox{ km s}^{-1}$ (e.g., Hobbs et al. 2005; Faucher-Gigu\`ere \& Kaspi 2006) and we would expect them to move in random directions. The tendency for the structures of ten PWNs to orient away from the center seems to suggest that the Galactic wind has a speed comparable to or greater than $\sim 400\hbox{ km s}^{-1}$. In the above scenario, the particle flow direction of the X-ray filament F3 should be from the southwest (closer to Sgr A*) to the northeast. This suggests that the pulsar, the origin site of the particles, is actually in the ``tail'' region defined in Fig 2. Then the evident spectral steepening from the southwest to the northeast (see section 3) can be naturally expained. Therefore, the spectral evolution along F3 also supports the existence of a radial high velocity wind in the GC region. While the X-ray filaments may not completely overlap with their radio NTFs, their overall orientations are similar. This is supported by the four X-ray NTFs (including filament F10 in our analysis) that have radio counterparts: G359.54+0.18 overlays exactly on the northern part of the two radio filaments (e.g., Yusef et al. 2005; Lu et al. 2003); G359.89-0.08 and its radio counterpart SgrA-E overlaps partly and extends in the same direction, with a centroid offset of $\sim 10 \arcsec$ (e.g., Yusef et al. 2005; Lu et al. 2003); G359.90-0.06 (SgrA-F) (e.g., Yusef et al. 2005) and G0.029-0.06 (F10) (see our subsection 3.10) also show similar spatial property. Generally speaking, the X-ray NTFs tend to be shorter than the radio NTFs. This centroid offset and smaller extent of X-ray filaments could be both explained by the much shorter synchrotron cooling lifetime in X-ray than in radio (e.g., Ginzburg \& Syrovatskii 1965). | 7 | 10 | 0710.2668 |
0710 | 0710.1946_arXiv.txt | We report on the first wide-field, very long baseline interferometry (VLBI) survey at 90 cm. The survey area consists of two overlapping 28 deg$^{2}$ fields centred on the quasar J0226$+$3421 and the gravitational lens B0218$+$357. A total of 618 sources were targeted in these fields, based on identifications from Westerbork Northern Sky Survey (WENSS) data. Of these sources, 272 had flux densities that, if unresolved, would fall above the sensitivity limit of the VLBI observations. A total of 27 sources were detected as far as $2\arcdeg$ from the phase centre. The results of the survey suggest that at least $10\%$ of moderately faint (S$\sim100$ mJy) sources found at 90 cm contain compact components smaller than $\sim0.1$ to $0.3$ arcsec and stronger than $10\%$ of their total flux densities. A $\sim90$ mJy source was detected in the VLBI data that was not seen in the WENSS and NRAO VLA Sky Survey (NVSS) data and may be a transient or highly variable source that has been serendipitously detected. This survey is the first systematic (and non-biased), deep, high-resolution survey of the low-frequency radio sky. It is also the widest field of view VLBI survey with a single pointing to date, exceeding the total survey area of previous higher frequency surveys by two orders of magnitude. These initial results suggest that new low frequency telescopes, such as LOFAR, should detect many compact radio sources and that plans to extend these arrays to baselines of several thousand kilometres are warranted. | \label{sec:introduction} The general properties of the 90 cm sky are not very well known and even less is known at VLBI resolution. Previous snapshot surveys at these wavelengths have only targeted the brightest sources and were plagued by poor sensitivity, radio interference and limited coherence times. Furthermore, the field of view that could be imaged was typically limited by the poor spectral and temporal resolution of early generation hardware correlators and the available data storage and computing performance at the time. As a result, although several hundred 90 cm VLBI observations have been made over the past two decades, images of only a few tens of sources have been published e.g. \citet{alt95}; \citet{laz98}; \citet{chu99}; \citet{cai02}. With such a small sample it is difficult to quantify the total population and nature of these sources. In particular, the sub-arcsecond and sub-Jansky population of 90 cm sources is largely unexplored. Recent improvements to the EVN hardware correlator at JIVE \citep{van04}, have enabled significantly finer temporal and spectral resolution. Combined with vast improvements in storage and computing facilities, it is now possible to image fields as wide as, or even wider than, the FWHM of the primary beam of the observing instrument. To complement the hardware improvements, new approaches to calibration and imaging have been developed to better utilise the available data and processing platforms. For example, \cite{gar05} performed a deep VLBI survey at 20 cm of a $36\arcmin$ wide field by using a central bright source as an in-beam calibrator. The approach was ideal for survey work as it permitted the imaging of many potential target sources simultaneously by taking advantage of the full sensitivity of the observation across the entire field of view. We have applied a similar technique at 90 cm by piggybacking on an existing VLBI observation of the gravitational lens B0218$+$357 and the nearby quasar J0226$+$3421, with the aim of surveying a 28 deg$^{2}$ field around each of the sources. The results provide an important indication of what may be seen by future low-frequency instruments such as the Low Frequency Array (LOFAR), European LOFAR (E-LOFAR) and the Square Kilometre Array (SKA). In this paper, we present the results of a 90 cm wide-field VLBI survey that covers two partially overlapping regions of 28 deg$^{2}$ each, surveying 618 radio source targets at angular resolutions ranging between 30 and 300 mas. For sources located at a redshift of $z=1$, the linear resolution corresponding to 30 mas is 230 pc. A \emph{WMAP} cosmology \citep{spe03} with a flat Universe, $H_{0}=72$ km s$^{-1}$ Mpc$^{-1}$ and $\Omega_{m}=0.29$ is assumed throughout this paper. | Our survey results indicate that at least $10\%$ of moderately faint (S$\sim100$ mJy) sources found at 90 cm contain compact components smaller than $\sim0.1$ to $0.3$ arcsec and stronger than $10\%$ of their total flux densities. This is a strict lower limit as the sensitivity of our observation was limited by the primary beam at the edge of the survey fields. None of the surveyed sources that were even slightly resolved by WENSS were detected. Similarly, none of the WENSS sources that were below the sensitivity limits of the VLBI observation were detected either, suggesting that none of these sources had significantly increased in brightness since the WENSS observations were carried out. The apparent lack of sources varying above our detection threshold must at least in part be due to resolution effects. As 90\% of the WENSS sources above the VLBI detection threshold are not detected they must be at least partially resolved at the VLBI resolution and the compact component of the radio emission would only be a fraction of the WENSS flux density. For the compact component of these sources to vary enough to be detected with VLBI, they must increase in strength by factors of perhaps at least a few (the reciprocal of the ratio of compact flux to WENSS flux) to be detectable with VLBI; the compact component of the flux needs to increase above the VLBI sensitivity limit. Resolution effects are masking variability in these sources. As discussed below, the detection of one apparent highly variable source in the VLBI data is rather remarkable. The interpretation of our detection statistics is complicated, in that the survey has a non-uniform sensitivity over both fields, due to the primary beam response of the VLBA antennas and the fact that we are imaging objects well beyond the half-power points of the primary beam. In addition, due to time and bandwidth smearing effects, as one images objects further from the phase centre, data on the long baselines is discarded, since the smearing effects make imaging difficult. A consequence of this is that the angular resolution is also non-uniform across the surveyed fields, with low resolution far from the phase centre. Not only is the flux limit variable across the field, the brightness temperature sensitivity also varies. It is possible to estimate the detection statistics of our survey for a uniform flux density and brightness temperature limit by considering sources not too far from the phase centre and for a flux sensitivity between the extremes at the phase centre and field edge. For example, if a sensitivity limit of 30 mJy beam$^{-1}$ is considered (achieved in the $0.25 - 0.5$ degree annulus of field 1 and in the $0.5 - 1$ degree annulus of field 2, and exceeded in the lower radius annuli in each field), 11 out of 55 possible sources are detected, a detection rate of 20\%, higher than the strict lower limit of 10\% estimated above for all sources at all annuli. \citet{gar05} performed a similar survey of the NOAO Bootes field at 1.4 GHz, using the NRAO VLBA and 100 m Green Bank Telescope. The survey covered a total of 0.28 deg$^{2}$, one hundredth of the area covered by our survey, and detected a total of 9 sources. The survey achieved sensitivities of $0.074-1.2$ mJy beam$^{-1}$ that enabled the detection of both weak and extended sources, whereas our 90 cm observations detected mainly compact sources or slightly resolved bright sources. Nonetheless, we can estimate the number of detections in this region that could be achieved using the 90 cm survey techniques described in this paper. The 0.28 deg$^2$ NOAO Bootes field contains a total of 13 WENSS sources, 6 of which have integrated flux densities $>30$ mJy. Based on our detection rate of 20\% for such sources we would expect to detect one WENSS source at 90 cm. Assuming a median spectral index of -0.77, only two of the \citet{gar05} sources have integrated flux densities above our 30 mJy beam$^{-1}$ limit at 90 cm, however, one of these is extended and would have a VLBI peak flux density that falls below our limit. Thus the observations of \citet{gar05} are consistent with our 90 cm VLBI results for sources with a peak flux density above 30 mJy beam$^{-1}$. Estimates of the percentage of sources detected with VLBI gives an estimate of the relative contribution of AGN (that contain compact radio emission and are detectable with VLBI) and starburst galaxies (which contain low brightness temperature radio emission not detectable with VLBI). Analysis of the ratio of starburst galaxies to AGN as a function of redshift (at high redshifts) can help to determine the initial sources of ionising radiation early in the Universe. As very little redshift data for our surveyed sources are available, such an analysis is not currently possible with this dataset. In practice, VLBI data at an additional frequency is also required, to confirm that the compact radio emission attributed to AGN has plausible spectral indices. The distribution of morphologies in the detected survey sources are typical of AGN. 10/27 sources are unresolved point sources, consistent with core-dominated AGN. A further 8/27 are clearly resolved into double component sources, consistent with being core-jet AGN or double-lobed radio galaxies. 7/27 sources have complex or extended structures, not obviously clear double components. Again, these sources may be core-jet AGN or radio galaxies. The remaining 2/27 sources are the gravitational lens and the quasar at the phase centres of the two fields. The serendipitous detection of a likely highly variable, very compact source near the target WENSS source B0223.8$+$3533 is intriguing. The total area imaged by this survey represents $\sim0.5\%$ of the area within the $0\arcdeg-2\arcdeg$ annulus and is equivalent to $\sim2.2\%$ of the FWHM of the VLBA primary beam. While it is difficult to place any limits on the real population of variable sources based on this one observation, it does highlight the importance of imaging wide-fields completely, in order to improve our understanding of such sources. \subsection{Future Prospects} The observations presented here demonstrate that extremely wide-field surveys can now be piggybacked on current and future VLBI observations at 90 cm. While this survey has mainly concentrated on detecting and imaging sources already detected by other surveys, we find tantalising evidence of a transient or highly variable source. We were fortunate to have found one that appeared in close proximity to one of our target sources but this may not always be the case. This provides a motive to take on a more ambitious survey of the entire field. Such a survey is not beyond the reach of current technology, it would require at most $\sim45$ times more processing compared to the project presented here, in order to image the entire primary beam of the VLBA using a similar faceted approach. While this is not the most efficient means of detecting transients, it will help progress the development of algorithms and techniques needed for next generation, survey-class instruments that operate at wavelengths or sensitivities not matched by current instruments. The observations presented in this paper were limited by the spectral and temporal resolution of the EVN correlator at the time of the observation. To minimise the effects of bandwidth and time-averaging smearing it was necessary to compromise resolution and image noise. Future technical developments in the capabilities of correlators will allow wide-field, global VLBI studies to be conducted without such restrictions. In particular, software correlators can provide extremely high temporal and spectral resolution, limited only by the time it takes to process the data \citep{del07}. They also allow for some pre-processing to be applied during the correlation process to, for example, mitigate the effects of radio interference or to correlate against multiple phase centres simultaneously. \subsection{Implications for LOFAR and SKA} The results of these observations provide important information on the nature and incidence of compact, low-frequency radio sources, with consequences for next generation, low-frequency instruments such as LOFAR and the SKA. LOFAR is currently being deployed across The Netherlands but remote stations are already under construction in neighbouring countries, in particular Germany. Other countries (e.g. UK, France, Sweden, Italy and Poland) are also expected to join this European expansion of LOFAR (E-LOFAR), extending the longest baseline from a few hundred, to a few thousand km. This development will provide LOFAR with sub-arcsecond resolution at its highest observing frequency (the $120-240$ MHz high-band). One concern associated with extending LOFAR to much longer baselines is whether enough cosmic sources will remain unresolved - this characteristic is required in order to ensure there are enough calibrator sources in the sky in order to calibrate the instrument across its full, very wide, field-of-view. The observations presented here suggest that at least one tenth of all radio sources (at the several tens of mJy level) are likely to exhibit compact VLBI radio structure in the LOFAR high-band. In all likelyhood, an even larger fraction of the E-LOFAR source population will therefore be bright and compact enough to form a grid of calibrator sources across the sky. From our results, we estimate the spatial number density of relatively bright (S$>10$ mJy) and compact (LAS$<200$ mas) sources at 240 MHz to be $\sim3$ deg$^{-2}$. The aggregate total of these compact sources within a beam should serve as a good calibrator for E-LOFAR and enable most of the low-frequency radio sky to be imaged with excellent sub-arcsecond resolution and high dynamic range. Extrapolation to LOFAR's low-band (10$-$80 MHz) is probably very dangerous, but there is every reason to believe that a large number of these sources will remain compact. In order to assess the relative numbers of starburst galaxies and AGN as a function of redshift, obviously large redshift surveys need to take place for these radio continuum objects. Such a survey could be conducted using the redshifted HI signal from these galaxies, using the SKA. | 7 | 10 | 0710.1946 |
0710 | 0710.1207_arXiv.txt | The total mass of a distant star cluster is often derived from the virial theorem, using line-of-sight velocity dispersion measurements and half-light radii, under the implicit assumption that all stars are single (although it is {\em known} that most stars form part of binary systems). The components of binary stars exhibit orbital motion, which increases the measured velocity dispersion, resulting in a dynamical mass overestimation. In these proceedings we quantify the effect of neglecting the binary population on the derivation of the dynamical mass of a star cluster. We find that the presence of binaries plays an important role for clusters with total mass $M_{\rm cl} \leqslant 10^5~{\rm M}_\odot$; the dynamical mass can be significantly overestimated (by a factor of two or more). For the more massive clusters, with $M_{\rm cl} \geqslant 10^5~{\rm M}_\odot$, binaries do not affect the dynamical mass estimation significantly, provided that the cluster is significantly compact (half-mass radius $\leqslant 5$~pc). | Young star clusters, with typical masses of $M_{\rm cl}=10^{3-6}~{\rm M}_\odot$, indicate recent or ongoing violent star formation, and are often triggered by mergers and close encounters between galaxies. Only a fraction of these young massive star clusters evolve into old globular clusters, while the majority ($60-90\%$) will dissolve into the field star population within about 30\,Myr (e.g., de Grijs \& Parmentier 2007). In order to understand the formation and fate of these clusters, it is important to study these in detail, and obtain good estimates of the mass, stellar content, dynamics, and binary population. The dynamical mass for a cluster in virial equilibrium, consisting of single, equal-mass stars, is given by: \begin{equation} M_{\rm dyn} = \eta \, \frac{ R_{\rm hm} \sigma_{\rm los}^2 }{G} \end{equation} (Spitzer 1987), where $R_{\rm hm}$ is the (projected) half-mass radius, $\sigma_{\rm los}$ the measured line-of-sight velocity dispersion, and $\eta \approx 9.75$. For unresolved clusters, $\sigma_{\rm los}$ is usually derived from spectral-line analysis, neglecting the presence of binaries. However, observations have shown that the majority of stars form in binary or multiple systems (e.g., Duquennoy \& Mayor 1991; Kouwenhoven et al. 2005, 2007; Kobulnicky et al. 2007). When binaries are present, $\sigma_{\rm los}$ does not only include the motion of the binaries (i.e., their centre-of-mass) in the cluster potential, but additionally the velocity component of the orbital motion. This results in an overestimation of the velocity dispersion, and hence of $M_{\rm dyn}$. | 7 | 10 | 0710.1207 |
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0710 | 0710.3202_arXiv.txt | We present spectroscopic observations of the quiescent black hole binary A0620-00 with the the 6.5-m Magellan Clay telescope at Las Campanas Observatory. We measure absorption-line radial velocities of the secondary and make the most precise determination to date ($K_{2}=435.4\pm0.5$ km s$^{-1}$). By fitting the rotational broadening of the secondary, we refine the mass ratio to $q=0.060\pm0.004$; these results, combined with the orbital period, imply a minimum mass for the compact object of $3.10\pm0.04$ M$_{\sun}.$ Although quiescence implies little accretion activity, we find that the disc contributes $56\pm7$ per cent of the light in B and V, and is subject to significant flickering. Doppler maps of the Balmer lines reveal bright emission from the gas stream-disc impact point and unusual crescent-shaped features. We also find that the disc centre of symmetry does not coincide with the predicted black hole velocity. By comparison with SPH simulations, we identify this source with an eccentric disc. With high S/N, we pursue modulation tomography of H$\alpha$ and find that the aforementioned bright regions are strongly modulated at the orbital period. We interpret this modulation in the context of disc precession, and discuss cases for the accretion disc evolution. | A0620-00 (V616 Mon) is the prototype Soft X-ray Transient, a class of low-mass binary stars which exhibit infrequent but intense X-ray bursts (Gelino, Harrison, and Orosz, 2001). In 1975 it became the brightest X-ray nova ever detected, at approximately 50 Crab \citep{Elvis75}, and it was the first nova to be identified with a black hole primary (McClintock \& Remillard, 1986; hereafter MR86). MR86 measured an orbital period of 7.75 hr and a radial velocity semiamplitude for the K-type secondary of 457 km s$^{-1}$, leading to a mass function $f(M)=3.18$ M$_{\sun};$ estimates of $K_{2}$ and $f(M)$ have decreased slightly since then (i.e. 433 km s$^{-1}$ and 3.09 M$_{\sun}$) (Marsh, Robinson, \& Wood 1994, hereafter MRW94). Given this minimum mass, it is likely that A0620-00 is a black hole. A substantial amount of work has gone into the analysis of A0620-00 in the last twenty years, with particular emphasis on ellipsoidal variations in the light curve and the contamination of the K-star flux by light from the accretion disc. As yet, no real consensus has been reached, mostly due to the complexity of the lightcurves. While ellipsoidal variations are obvious, they are highly asymmetric (Leibowitz, Hemar, \& Orio 1998); the origin of the asymmetry is undetermined. Modelling this lightcurve, \citet{Gelino01} determined in inclination of 41$\pm 3\degr$, invoking starspots to explain the asymmetries. Shahbaz, Naylor, and Charles (1994) found a 90 per cent confidence interval of $i=$30--45$\degr$ given the mass ratio of A0620-00, modelling their asymmetries with the bright spot where the accretion stream hits the disc. Lightcurve modelling is also complicated by the variability of the disc itself. To quantify ellipsoidal variations, most authors assume the disc to be constant, and justify the claim by noting that A0620-00 is quiescent. They do not mention that estimates of the disc contamination range from $<$3 per cent \citep{Gelino01} to $\la50$ per cent (MR86). The contribution from this disc is not only unclear, but apparently not constant. More than half of A0620-00's 58-year burst cycle has passed, and it is important to note that quiescent does not mean inactive. We will argue that the variability of the accretion disc cannot be neglected. In order to determine definitively the mass of the compact object, it is very important to understand the structure and variation of the accretion disc. MRW94 made enormous progress towards this goal. In 2004, Shahbaz et al. (hereafter S04) noticed signatures of an eccentric disc not seen in previous Doppler maps, but lacked the phase coverage to verify their hypothesis. Therefore, as follow-up to the work of MRW94 and S04, and as part of a Doppler imaging survey of black hole and neutron star binaries, we undertook phase-resolved optical spectroscopy of A0620-00. In $\S 2$ we describe our observational methods and data reduction. In $\S 3$ we measure the radial velocity of the secondary star, the system mass ratio, and attempt to quantify flickering. In $\S 4$ we present Doppler images of the accretion disc at several wavelengths, investigate evidence for an eccentric disc, and report results of modulation tomography of the H$\alpha$ line. We discuss conclusions from the variability of the disc and our Doppler maps in $\S 5.$ | We have presented spectroscopic analysis of the black hole binary A0620-00. We measure an absorption-line radial velocity $K_{2}=435.4\pm0.5$ km s$^{-1}$. With two measurements of the rotational broadening of the secondary, we find a mass ratio of $q=0.060\pm0.004$ and a minimum mass of $3.10\pm0.04$ M$_{\sun}$ for the primary object. With the most likely inclination of 41$\degr$ from \citet{Gelino01}, measured in $J,~H,$ and $K$, the black hole has a mass of 11.1 M$_{\sun}.$ The strong infrared flickering discussed earlier, in conjunction with unexplained smooth variability in the lightcurve and uncertainty in the disc spectrum itself, makes it difficult to estimate the true uncertainty in the inclination. The range of reported inclinations, 31$\degr$ to 70.5$\degr$(Gelino et al. 2001 and references therein), results in a black hole mass between 3.7 M$_{\sun}$ and 22.7 M$_{\sun}$. Until the nature and variability of the light from the disc is revealed in full, this conservative error estimate must be sufficient. We also find that the secondary contributes $44\pm7$ per cent of the light near 5500 \AA. As this means that the disc contributes a significant fraction of the light, especially in emission line regions, it becomes important to assess the variability of the disc, particularly if the inclination is to be determined by lightcurve modelling. As noted, it is common to assume that the disc is a constant source of light. While it may be valid when $f$ is large, three observational points cast doubt on this assumption: \begin{enumerate} \item S04 performed a detailed study of flares from A0620-00, which in their observations have amplitudes nearing 20 per cent of the source flux. \item Measurements of the fraction $f$ of light contributed by the secondary have not been consistent. MR86 found $40\pm10$ per cent at 5100 \AA, MRW94 found $94\pm3$ per cent at H$\alpha$, and \citet{Gelino01} found $f\ga97$ per cent in $J,~H,$ and $K$, assuming a \textit{constant} diluting source of light. \item Observations in all four Spitzer bands, taken approximately two weeks before our observations on the Clay telescope, show strong flickering which is highly correlated with simultaneous R-band lightcurves from the 1.2m telescope on Mt. Hopkins. \end{enumerate} While McClintock, Horne, and Remillard (1995) rightly point out that the absorption line strength of the template will affect the observed dilution fraction, unless it can be shown that the measured fraction is strongly correlated with template line strength, a physical origin for this variation cannot be ruled out. Future studies could assess the true dependence of $f$ on template star, as well as the long-term variability of $f$, by observing both BS 753 (MWR94's template) and HD 7142. Our measurements of the H$\alpha$ equivalent widths, when compared to those of MRW94, suggest a physically real origin for the variation, because equivalent widths are independent of the template star and the instrument. So there appears to be evidence for a disc which is brighter relative to the secondary than it used to be. Indeed, the modulations of the equivalent width even point to a physically non-uniform flickering, because the disc emission lines vary relative to the non-stellar continuum. We can tentatively identify the flickering with the crescents, neither of which was present in 1994, so our conclusion seems viable. The line emission and the continuum may have different radial emissivity dependencies, which could result in slower modulations. We also detect several bright regions in the disc: one near the gas stream impact point, and two crescent-shaped regions on opposite sides of the disc. The reality of these features, as well as the non-uniform flickering, are confirmed by modulation tomography of the H$\alpha$ disc line, which reveals variation near the bright crescents. First noticed by S04, the crescent-like features may indicate an eccentric disc, which is predicted for systems like A0620-00 with small mass ratios and large discs. It seems that we have observed an eccentric disc, but let us consider the evidence. From our observations, the following are clear: \begin{enumerate} \item The disc is bright and variable. The brightness is evident in the increased dilution of the secondary spectrum, and the variability is clear from a number of phenomena. First, the trailed H$\alpha$ line shows clear evidence of flickering events. Second, the subtracted equivalent width of the same line is variable beyond explanation by noise alone. Third, the phase-resolved light fraction cannot be reproduced by ellipsoidal variability on top of a constant source of light. \item The disc extends to the 3:1 tidal resonance. This is a simple point, clear from the Doppler and modulation maps. \item The disc is not centered on the radial velocity of the black hole. \item The disc is not radially symmetric, but characterized by bright crescent-shaped regions. \item The crescent regions are modulated at the orbital period. \end{enumerate} Each of these points alone would be insufficient evidence to conclude that the accretion disc is eccentric and precessing. But with the exception of a direct image of the elliptical disc, we can present a complete and coherent argument that this is the case. The disc has grown to tidal resonance, where the enhanced disc viscosity results in bright and variable rims of extra dissipation. We then observe crescents of extra dissipation at relatively low velocities, as expected. Given the viscous effects, it is predicted that the disc will receive a gravitational torque from the secondary, and begin to precess. The asymmetries introduced here shift the velocity center of disc emission away from the black hole, and we find that the disc is not centered on the black hole in velocity space. Furthermore, given the beat period between the precession and orbital motion, the regions of viscous dissipation should be modulated at roughly the orbital period. Modulation tomography reveals this to be the case. In retrospect, knowing that portions of the disc are modulated on the orbital period, we look closer at the fraction of light contributed by the secondary star, and see that it is not well fit by ellipsoidal modulations for a system at the inclination of A0620. But if another component of the system was variable on the orbital period, as we have observed the disc to be, then there is no need for concern. The physical picture, a precessing elliptical disc torqued by the secondary star, predicts and produces all the phenomena we have discussed in our data, which are of high quality. To put it another way, the eccentric disc hypothesis is nicely self-consistent. It explains why and how the accretion disc has changed, allows the disc to be large enough for the growth of eccentric modes, and predicts the phenomena that we observe. It is unfortunately not possible at this point to make an estimate of the disc eccentricity. \citet{Smith07} have shed a great deal of light on the evolution of disc eccentricity and energy dissipation with 3D SPH simulations. They find that systems with $q$ between 0.08 and 0.24 develop low-mass eccentric discs withsuperhumps; for $q=0.0526,$ the disc exhibits a short-lived superhump and decaying eccentricity. All mass ratios show enhanced dissipation in the disc from the thermal-tidal instability, even without the eccentric modes. Since we have not observed a superhump, we cannot place A0620-00 in either category. If it falls in the more extreme group, the disc eccentricity is likely zero (reached after about 300 orbital periods) \citep{Smith07}. In that case, the steady state is a massive disc. If the steady state is very long-lived, and the disc continues to grow, this could explain the enormous intensity of novae like A0620-00. If it fits among the less extreme mass ratios, the disc eccentricity is around 0.1--0.2, and a superhump should be observable with better photometry and a longer baseline \citep{Smith07}. A0620-00 may also be at a transition between those cases, and its evolution might be somewhat more erratic, as suggested by the SMARTS data discussed earlier. For example, it may toggle between states of quiescence, superhumps, and variability (like what we have observed here). It might, then, be erroneous to interpret this recent increase in brightness as the build towards outburst. While we have strong evidence that the accretion disc around the black hole has grown out to the tidal distortion radius, evolved into an eccentric disc, and started to precess, further study is required to verify our conclusion. Data from SMARTS, FLWO, and Spitzer will further quantify flickering, and may reveal a superhump, or some new period consistent with our results, and future programs of tomography will track the evolution of the accretion disc. In anticipation of the impending outburst, and in light of progress in simulations, we suggest that this well-studied system not be disregarded or ignored, for it affords us the opportunity to watch the evolution of an accretion disc from quiescence to outburst, and the chance to test models for disc instabilities in X-ray novae. | 7 | 10 | 0710.3202 |
0710 | 0710.3805_arXiv.txt | Multi-GeV and TeVs gamma sources are currently observed by their Cherenkov flashes on Telescopes (as Magic, Hess and Veritas), looking vertically up into sky. These detectors while pointing horizontally should reveal also the fluorescence flare tails of nearby down-going airshowers. Such airshowers, born at higher (tens km) altitudes, are growing and extending up to lowest atmospheres (EeVs) or up to higher (few km) quotas (PeVs). These fluorescence signals extend the Cherenkov telescopes to a much higher Cosmic Ray Spectroscopy. Viceversa, as it has been foreseen \cite{Fargion2005} and only recently observed, the opposite takes place. Fluorescence Telescopes made for UHECR detection (as AUGER ones) may be blazed by inclined Cherenkov lights: less energetic, but frequent (PeVs) CR are expected to be often detected. Nearly dozens blazing Cherenkov at EeV should be already found each year in AUGER, possibly in hybrid mode (FD-SD, Fluorescence and/or Surface Detector). Many more CR events (tens or hundred of thousands) at PeVs energies should be blaze Cherenkov lights each year on the AUGER Fluorescence Telescopes. Their UV filter may partially hide their signals and they cannot, unfortunately, be seen yet in any hybrid mode. At these comparable energy the rarest UHE resonant antineutrino $\bar{\nu}_e+e$ interactions in air at $\frac{{M_{W}}^2}{2m_e}=6.3$ PeV energy, offer enhanced $W^-$ Neutrino Astronomy showering at air horizon, at $\sim90^\circ$, while crossing deep atmosphere column depth or Earth (Ande) boundaries. However, AUGER FD are facing opposite way. An additional decay channel rises also (after resonant neutrino skimming Earth) via their secondary $\tau$ exit in air, by decay in flight via amplified showering: $\bar{\nu}_e+e\rightarrow W^-\rightarrow\bar{\nu}_{\tau}+\tau$. Moreover, expected horizontal UHE GZK neutrinos $\nu_{\tau}\,\bar{\nu}_{\tau}$ at EeVs energy, powered by guaranteed cosmogenic GZK \cite{Greisen:1966jv, za66}, $\nu_{\mu}\,\bar{\nu}_{\mu}$ flavor conversions (in cosmic distances), are also producing penetrating UHE EeV lepton taus that could sample, better and deeper than PeVs ones, the Earth skin. Such almost horizontal and up going tau showers, originated by UHE astronomical neutrino, may shower and flash by Fluorescence and/or Cherenkov diffused lights at Auger Sky in a few years (nearly three). Viceversa, at Hess, MAGIC and VERITAS Horizons, at tens or a hundred kilometer distances, the same up going $\tau\,\bar{\tau}$ airshowers might rise via fluorescence. On axis they might blaze (rarely) as a Cherenkov flashes below the horizons, possibly correlated to BL Lac or GRB activity. Also UHE ($1-0.1$ EeV) GZK $\tau$ showering, can be observed upward once reflected onto clouds. The geomagnetic splitting may tag the energy as well as the inclined shower footprint as seen in a recent peculiar event in AUGER. Additional stereoscopic detection may define the event origination distance and its consequent primary composition, extending our understanding on UHECR composition. | \begin{figure} \centering \includegraphics[width=7.5cm]{Fargion_Ricap1.eps} \caption{The different geometry for a Magic-like telescope searching at horizon airshower tails. Nearby PeVs CR can be observed by fluorescence at few km altitude while a more powerful airshower (EeV), might be observed at lower altitudes because of the larger slant depth on far edges. Cherenkov airshowers might blaze in the telescope if in axis. Rarest up-going Tau airshower may escape the Earth at far horizon, leading to up-going flashes. Guaranteed nearby CR (PeVs-EeVs) airshower Cherenkov lights, may also be observed by reflection from nearby hills or mountains (as in Magic or Veritas), or from the sea \cite{Fargion07}}\label{fig1} \end{figure} \label{intro} Cosmic Rays is a mature, century old Science. It still hides its secrets beyond the amazing homogeneity and isotropy in all energy range. Only photons, our best neutral courier in Astronomy, offered up to now a view of the Universe from lowest (radio) to highest (TeVs) energies. Because of the TeVs-PeVs-EeVs photon self-interaction with relic IR (Infrared), CBBR (Cosmic Black Body Radiations) and Radio backgrounds, UHE (Ultra High Energy) photons are bounded in nearby Universe. Therefore neutral neutrinos may offer a far insight of CR sources because of their weak interaction for we could also reveal the inner core of their mysterious accelerators. After a decade, the exciting hopes of a discover by AGASA and by Hires of (an almost undeflected) UHECR traces, a long waited new Particle Astronomy have been frustrated by AUGER results. No UHECR clusterings toward BL Lacs or AGN seems to arise up now. Even if GZK cut-off (the CBBR opacity to nucleons above few $10^{19}$ eV flux \cite{Greisen:1966jv, za66}) appeared to be finally confirmed by Hires \cite{Hires06} and AUGER \cite{Yamamoto2007}: no Galactic or nearby Universe map within GZK cut off volume seems to be correlated with these UHECR events. If homogeneity and isotropy will survive AUGER, Z-Burst model \cite{Fargion-Mele-Salis99}, linking far cosmic UHE ZeV $\nu$ sources scattering on $\bar{\nu}$ relic ones, remains the unique natural option. Otherwise, our Near Universe as Virgo and our Super- Galactic group and/or plane, must rise soon. Moreover an unexpected heavy composition of extreme UHECR makes more urgent an independent UHECR spectroscopy. As well as the discover of UHE $\nu\,\bar{\nu}$ GZK, rare but guaranteed GZK \cite{Greisen:1966jv, za66} secondary neutrino traces. To this project and to its solution we address in present paper. \begin{figure}[] \centering \includegraphics[width=7.5cm]{Fargion_Ricap2.eps} \caption{The rare inclined UHECR event seen in axis from above. To this picture derived by an AUGER presentation, we overlapped a drawing of the airshower components, their bending and the consequent geomagnetic $\vec{B}_\oplus$ splitting. Note the $\vec{B}_\oplus$ vector pointing to North and upward respect the ground. Note also the consequent vertical and lateral charge bundle separations. Each charge-bundle follow its bend trajectory that generate its own Cherenkov beam. The comparability between $\vec{B}_{\oplus\perp}$ and $\vec{B}_{\oplus\parallel}$ field vector modules, is the cause of a similar angular (lateral-vertical) deflection. Their projection on the ground is not symmetric at all. The dashed-ellipse on the right side marks our forecast Cherenkov spot made by $e^+$ split shower component, undetected (out of a Cherenkov reflection on the ground hardly recorded) by Los Leones FD station which instead detected the main Florescence flare. The top-left side ellipse, marks the probable Cherenkov spot born by negative electrons blazing on Coihueco telescope.}\label{fig2a} \includegraphics[width=7.5cm]{Fargion_Ricap3.eps} \caption{The same Auger event seen from another angulation \cite{Auger07}. The inclined UHECR is reaching the ground on the center, where the muon pairs are clustering into a twin overlapped ellipses. The arrival angle is approximated at $80^\circ$ zenith angle \cite{Auger07}. Our consequent estimated primary energy is above few EeV (possibly around ten EeV). The Cherenkov blaze time on Coihueco $\tau'$, being slightly off axis ($\sim5^\circ$), must appear much shorter (one or two order of magnitude) than the Fluorescence duration signal $\simeq\tau$ reaching Los Leones, because of both the geometric and the relativistic shrinkage, $\tau'=\tau_o\cdot(1-\beta\cos(\theta))$. The FD, because of lack of angular resolution in Auger Telescope, might be unable to reveal the electron pair splitting. The Cherenkov bending angle extends a few degrees: at Coihueco the airshower beam is roughly at $5^\circ$ above horizons, while its shower beginning reaches $7^\circ-9^\circ$ angle. Therefore the Cherenkov splitting shape might be marginally detectable by Coihueco Telescope, by a two-three pixels separation along its inclined polarized axis.}\label{fig2b} \end{figure} \subsection{Fluorescence flares within Cherenkov Telescope} We foresee that Cherenkov telescopes while pointing at horizons ($\geq80^\circ$ zenith angles) may observe a truncate image (a cylinder like) of downward fluorescence airshower, lightening (See Fig.\ref{fig1}). These flaring views may appear often at few tens km distance, toward $80^\circ$ angle for tens $PeVs$ (or a hundred km at $\simeq85^\circ$ for EeVs energies), as often as once a night. Nearby hills or reflecting sea may disturb the detection. We foresee that such a discover must occur soon, amplifying MAGIC, VERITAS and HESS high energy CR yields. \subsection{Cherenkov blazing photons on Fluorescence Telescopes} The opposite also take place: Cherenkov photons may hit Fluorescence Telescopes, even if most Fluorescence detectors are masked by UV filter. Indeed the blazing Cherenkov lights are collimated into a narrow cone; therefore they are more rarer than Fluorescence isotropic signals. We may estimate that few hundreds of EeV airshowers in a year might hit with Cherenkov the AUGER FD. Probably only a few dozens are well revealed with the FD and by muons on SD, as well as Cherenkov lights, as it has been foreseen \cite{Fargion2005}. A very peculiar and pedagogical event has been shared on line by AUGER collaboration \cite{Auger07}. We shall analyze that event in the next sections. | Cherenkov and Fluorescence Telescope may enlarge their view and role tracing both lights. The wider CR range will allow a better spectroscopy at PeV-EeV (knee-ankle) regions and a more detailed anatomy of UHECR composition. MAGIC, HESS and VERITAS may soon trace the Fluorescence lights of downward UHECR airshowers. MAGIC and VERITAS must reveal the Cherenkov reflections also on nearby Mountains. The recent inclined UHECR event in Auger \cite{Auger07} clearly foreseen in \cite{Fargion2005} discussed in this paper offer the first footprint for such rich information derived by muons bundles, electron pairs showering and splitting into polarized Cherenkov and Fluorescence traces. Within this novel spectroscopy a hidden Neutrino Astronomy wait to be finally unveiled \cite{Fargion2007}. | 7 | 10 | 0710.3805 |
0710 | 0710.0421_arXiv.txt | {In this note we present global string solutions which are a generalization of the usual field theory global vortices when the kinetic term is DBI. Such vortices can result from the spontaneous symmetry breaking in the potential felt by a D$3$-brane. In a previous paper (hep-th/$0706.0485$), the DBI instanton solution was constructed which develops a "wrinkle" for stringy heights of the potential. A similar effect is also seen for the DBI vortex solution. The wrinkle develops for stringy heights of the potential. One recovers the usual field theory global string for substringy potentials. As an example of the symmetry breaking, we consider a mobile D$3$-brane on the warped deformed conifold. Symmetry breaking can occur if the structure of the vacuum manifold of the potential for the D$3$-brane changes as it moves through the throat region. } \begin{document} | String theory embeddings of inflationary scenarios typically involve D-branes which are either static or undergoing some motion in the transverse directions. The world volume theory of the D-brane is described by the DBI action. Since the DBI action is a higher dimensional generalization of the Lorentz invariant relativistic action for a point particle, relativistic effects are to be expected for the D-brane dynamics under appropriate relativistic conditions. Such examples of DBI relativistic effect are studied in \cite{Brown:2007ce,Brown:2007vh,Silverstein:2003hf, Chen:2004gc}. In \cite{Silverstein:2003hf}, the speed limit arising due to the relativistic nature of the DBI action is responsible for inflation. In \cite{Chen:2004gc} a variant of this scenario was proposed, again based on the DBI action. Further, the DBI effect has been shown to prevent the occurence of slow roll eternal inflation in DBI inflation scenarios \cite{Chen:2006hs}. These inflationary scenarios employ mobile branes and, therefore, when the D$3$-brane speed becomes relativistic see DBI effects. However, the nonlinear nature of the DBI action can lead to new behavior in certain phenomena even when the D-brane is static. In \cite{Brown:2007ce} the tunneling of a D$3$-brane from a metastable to a true vacuum was investigated along the lines of the usual quantum field theory analysis given in \cite{Coleman:1977py}. For substringy barriers the tunneling picture for the DBI action matches with the usual QFT picture given in \cite{Coleman:1977py}. But once the barrier between the metastable and the true vacuum assumes stringy heights, tunneling is much more enhanced than what one would have expected based on the usual QFT intuition. This is surprising as the usual QFT intuition has taught us to expect an exponential suppression in the tunneling probability with an increase in the height of the barrier. This counterintuitive result is easily understood, however, based on a similar treatement of a point charged particle tunneling through an electrostatic barrier \cite{Brown:2007vh}. Once the height of the barrier is large enough, there is a significant Schwinger pair production of point particle- antiparticle and this Schwinger effect can enhance tunneling. Similar effect for the D$3$-brane enhances the tunneling rate. The D-brane, therefore, need not be moving at relativistic speeds to see interesting DBI effects. In this paper we investigate global vortex solutions in the world volume of a D$3$-brane described by the DBI action. Similar to the findings in \cite{Brown:2007ce}, these vortex solutions display a departure from the usual field theory behavior once the height of the potential of the complex scalar field becomes stringy. The DBI vortex solutions that we will construct should not be confused with the vortex solutions of the tachyon field formed due to tachyon condensation after brane-antibrane annihilation (constructed, for example, in \cite{Jones:2002si} ) which correspond to codimension two branes. The DBI vortices we construct are generalizations of the well known field theory global vortices when the kinetic term is DBI. Indeed, for small gradients (which will correspond to small heights of the potential) the DBI vortex resembles the usual field theory vortex. Further, these vortices can form even when the separation between a brane and an antibrane is much greater than the critical distance when tachyon condensation initiates. For the tachyon vortices to form, the brane-antibrane separation should be string length. | We have considered the DBI action for a complex field $\psi$ with a potential term $V(\psi)$ and constructed the vortex solutions which are generalizations of the usual global vortex of the complex scalar field theory. For substringy heights of the potential the vortex resembles the usual field theory global vortex. For stringy potentials the vortex develops a wrinkle analogous to the instanton wrinkle found in \cite{Brown:2007ce} \footnote{Vortex solutions for the DBI action have been studied by various authors including \cite{Callan:1997kz, Gibbons:1997xz, Hashimoto:1997px,Hashimoto:2003pu}. These vortex solutions (BIons) can develop a throat and have double valued $\phi$.}. We have neglected the world volume gauge fields. Including the world volume gauge fields might lead to the DBI analogues of field theory local strings. These vortices should not be confused with the vortex solutions on a tachyon profile that are produced upon tachyon condensation initiated by the brane-antibrane instability. Unlike tachyon vortices, the DBI vortices can form even when the brane-antibrane separation is large in string length units. A similar effect was seen in the context of DBI instanton in \cite{Brown:2007ce} where the instanton interpolating between the metastable and the true vacuum of the potential $V(\phi)$ was studied in the thin wall limit. When the height of the barrier $V_0 < 1$ one gets the usual QFT instanton. However when $V_0 > 1$ the instanton develops a wrinkle due to the double valuedness (and turning around) of the instanton configuration. For the vortex solution the wrinkle appears at $V_0 \approx 0.8$. This difference in the critical $V_0$ is due to the friction term present in the vortex equation of motion (Eq.(\ref{dbi})). The wrinkled instanton in \cite{Brown:2007ce} was constructed in the thin-wall limit where the friction term is set to zero. The appearance of the wrinkle can be understood by following the analogy with the DBI instanton of \cite{Brown:2007ce}. Assuming that the vortex configuration will always have a large enough value of the radial coordinate $r$ (i.e. apriori assuming the formation of a wrinkle), the second and the third terms in Eq.(\ref{dbics}) (i.e. $ \frac{\gamma}{r} \frac{d\chi}{dr}$ and $\gamma \frac{n^2\chi}{r^2} $) can be neglected. The equation of motion then resembles the thin wall equation of motion for the DBI instanton in \cite{Brown:2007ce} and following the reasoning for the instanton, a wrinkle must develop. This is, of course, a heuristic way to see the appearance of a wrinkle in the vortex, which will fail to give the exact value of the radial coordinate $r$ or height of the potential $V_0$ where the wrinkle appears. The asymptotics of the DBI vortex at radial infinity will remain the same as that for the usual field theory vortex of Sec.(\ref{qft}). This is because at radial infinity the gradient term vanishes and the DBI kinetic term does not play any role. The DBI effect only occurs near the origin (i.e. near the core of the vortex) where the gradient term is big and sources the DBI term. Since the DBI global vortex asymptotics at radial infinity are the same as the usual field theory global vortex asymptotics at infinity, there will still be a logarithmic divergence in the total energy of the field configuration. However, as explained before, inspite of the double-valuedness of the DBI vortex field configuration, the Hamiltonian density will remain finite everywhere. In Sec. (\ref{moduli}) we considered the possibility of the cosmological formation of these vortices in a brane inflation scenario. Brane inflation consists of two steps. First, before the inflation can begin, a D$3$-antibrane migrates down the throat and settles at the tip. In the presence of an appropriate moduli space, vortices can form on the antibrane world volume when it settles at the tip. However, in the next stage a D$3$-brane is attracted to the antibrane and this leads to inflation. The inflationary stage washes away the vortices. Further, the different domains on the antibrane which have the $\psi$ field at different points of the $S^1$ become exponentially large during inflation. Consequently, the D$3$-brane effectively falls towards a single point on the $S^1$. No DBI vortices form at the end of inflation. Our approach in this note has been phenomenological. Motivated by the results for the D$3$-brane moduli space on a warped deformed conifold in \cite{DeWolfe:2007hd}, we consider a potential $V(\psi)$ along the $S^3$ angular coordinates for the D$3$-brane and examine the global vortex solution. Whether or not a wrinkle will form depends on the height of the potential. If $V(\psi)$ never attains stringy heights, then there is no possibility of a wrinkle. At this point we simply note that the simplest DBI inflation scenario \cite{Silverstein:2003hf} considers a DBI action of the form Eq.(\ref{eom}) (with gravity added). The inflaton potential $V(\phi)$ is generated by radiative or bulk effects whose value must be large when measured in terms of the warp factor/local string units $f(\phi)$. In particular, power law inflation occurs at late times (i.e. when $t \to \infty$) as the mobile D$3$-brane is nearing the tip. The warp factor and the potential have late time dependence $f(\phi) \to t^4$ and $V(\phi) \to 1/t^2$ which leads to $f(\phi)V(\phi) \to t^2$ in the late time DBI regime. In such a setting it appears that one can generate large potentials needed for interesting DBI effects. For the appearance of a wrinkle we require a potential whose value is order one in local string units. Another issue while considering the appearance of the wrinkle is the trustability of the DBI description. The DBI description is valid whenever the extrinsic curvature of the solution $\psi(r)$ is low. In the absence of any warp factor this extrinsic curvature is given by \cite{Brown:2007ce, Sarangi:2007jb} \baray K(\psi) = \frac{1}{\sqrt{\alpha'}} \frac{\partial V(\psi)}{\partial \psi}, \earay i.e. as long as the slope of the potential is small in string units, the DBI action has small higher derivative corrections. It would be interesting to see if there is some version of brane inflation on a warped deformed conifold which can lead to the formation of these DBI vortices. Further, it would be interesting to find the local string version of these DBI global vortices as local strings can be realistic cosmic string candidates. If such local strings still develop a wrinkle due to the DBI effect, this wrinkle would be the result of the stringy DBI action. Such cosmic strings would, in principle, differ from their field theory cousins and perhaps lead to novel phenomenological consequences. Further one could construct DBI defects, that are extensions of the DBI vortices we have studied, on the world volume of multiple D$3$-branes. The theory would then be a non-abelian one. One could then look , for example, for the DBI extension of the t' Hooft-Polyakov monopole. It is tempting to speculate that just as what we saw for the DBI vortex solution, there will be no DBI effects present at radial infinity of the monopole solution. This is because the field gradient vanishes at large distances from the core of the monopole. The DBI effect would occur for large field gradients near the core of the monopole and lead to the formation of a wrinkle near the core for potentials with stringy heights. | 7 | 10 | 0710.0421 |
0710 | 0710.2338_arXiv.txt | We present the first large-scale effort of creating composite spectra of high-redshift type Ia supernovae (SNe~Ia) and comparing them to low-redshift counterparts. Through the ESSENCE project, we have obtained 107 spectra of 88 high-redshift SNe~Ia with excellent light-curve information. In addition, we have obtained 397 spectra of low-redshift SNe through a multiple-decade effort at Lick and Keck Observatories, and we have used 45 ultraviolet spectra obtained by \hstiue. The low-redshift spectra act as a control sample when comparing to the ESSENCE spectra. In all instances, the ESSENCE and Lick composite spectra appear very similar. The addition of galaxy light to the Lick composite spectra allows a nearly perfect match of the overall spectral-energy distribution with the ESSENCE composite spectra, indicating that the high-redshift SNe are more contaminated with host-galaxy light than their low-redshift counterparts. This is caused by observing objects at all redshifts with similar slit widths, which corresponds to different projected distances. After correcting for the galaxy-light contamination, subtle differences in the spectra remain. We have estimated the systematic errors when using current spectral templates for K-corrections to be \about 0.02 mag. The variance in the composite spectra give an estimate of the intrinsic variance in low-redshift maximum-light SN spectra of \about 3\% in the optical and growing toward the ultraviolet. The difference between the maximum-light low and high-redshift spectra constrain SN evolution between our samples to be $< 10$\% in the rest-frame optical. | \label{s:intro} Type Ia supernovae (SNe~Ia) are the most precise known distance indicators at cosmological redshifts. The meticulous measurement of several hundred SNe~Ia at both low and high redshifts has shown that the expansion of the Universe is currently accelerating \citep{Riess98:lambda, Riess07, Perlmutter99, Astier06, Wood-Vasey07}; see \citet{Filippenko05} for a recent review. The underlying assumption behind that work is that high-redshift SNe~Ia have the same peak luminosity as low-redshift SNe~Ia (after corrections based on light-curve shape; e.g., \citealt{Phillips93}). The luminosity of a given SN and its light-curve shape are determined by initial conditions of the white dwarf progenitor star (e.g., mass at explosion, C/O abundance, and metallicity), and the properties of the explosion (e.g., deflagration/detonation transition, the amount of unburned material, and the density at the ignition point). The progenitor properties are set by the initial conditions at the formation of the progenitor system, presumably having properties similar to the global galactic properties at that time. Since low-redshift SN~Ia progenitor systems likely form, on average, in significantly different environments than high-redshift SN~Ia progenitors, one may assume that some amount of evolution is inevitable \citep[for a discussion of different causes and effects of SN evolution, see][]{Leibundgut01}. Theoretical studies of SN~Ia evolution have focused on the composition, particularly metallicity, of the progenitor system as the primary potential difference between the two samples. There have been two major studies with conflicting results. For their study, \citet{Hoflich98} changed the progenitor metallicity and modeled the explosion, including a full nuclear-reaction network. \citet{Lentz00} changed the metallicity of the results of W7 models \citep{Nomoto84} and input those parameters into their PHOENIX code \citep{Hauschildt96} to produce synthetic spectra. The main difference between these methods is the definition of ``metallicity.'' \citet{Hoflich98} uses the term to mean the metallicity of the progenitor star, while \citet{Lentz00} uses it to mean the metallicity of the ejecta. The differing definitions of metallicity yield different initial conditions, which resulted in contradictory results from these studies. \citet{Hoflich98} suggest that with increasing metallicity, the ultraviolet (UV) continuum of the SN increases, while \citet{Lentz00} suggest that it decreases. Ultimately, the differences are the result of differing density structures \citep{Lentz00, Dominguez01}. Although the method of \citet{Lentz00} seems less physical than that of \citet{Hoflich98} (simply scaling the metallicity of the ejecta by solar abundances does not take into account, for example, that the Fe-group elements are mainly produced in the SN explosion), they provide model spectra for varying metallicities, which may elucidate differences between low and high-redshift SN spectra. Further predictions for lower metallicity include faster rise times \citep{Hoflich98}, faster light-curve decline \citep{Hoflich98}, lower $^{54}$Fe production \citep{Hoflich98}, smaller blueshifting of \ion{Si}{2} $\lambda 6355$ \citep{Lentz00}, decreasing $B-V$ color \citep{Dominguez01, Podsiadlowski06}, and changing luminosity \citep{Hoflich98, Dominguez01, Podsiadlowski06, Timmes03}. \citet{Ropke04} suggest that the C/O ratio of the progenitor does not significantly affect peak luminosity. Observationally, a lack of evolution has been supported by investigating various SN quantities such as rise time \citep{Riess99:risetime}, line velocities \citep{Blondin06, Garavini07}, multi-epoch temporal evolution \citep{Foley05}, line strengths \citep{Garavini07}, and line-strength ratios \citep{Altavilla06}. There have also been studies comparing the spectra of individual high-redshift SNe~Ia to low-redshift SNe~Ia \citep{Riess98:lambda, Coil00, Hook05, Matheson05, Balland07}, all of which have concluded that there is no clear difference in spectral properties between the two samples. \citet{Bronder07} recently presented measurements of line strengths that suggest a difference between low and high-redshift SNe~Ia in one of three features measured. They find that the difference is highly dependent on the galaxy contamination at high redshift and might be affected by their small low-redshift SN sample. Consequently, they note that the difference is interesting but not significant. Despite the consistencies in spectral properties, \citet{Howell07} note a slight shift in the mean photometric properties of SNe~Ia with redshift. They explain this evolution as a change in the ratio of progenitors from the ``prompt'' and ``delayed'' channels \citep{Scannapieco05}, corresponding to young and old progenitor systems at the time of explosion, respectively. In particular, the light-curve shape parameter ``stretch'' \citep{Goldhaber01} increases with redshift. \hst\, observations of ESSENCE objects suggested that the sample may have a large proportion of objects with slow-declining (large stretch) light curves, but this is probably the result of a selection bias\citep{Krisciunas05}. Since stretch (and other luminosity light-curve parameters) is correlated with spectral properties, one might expect the spectra of high-redshift SNe~Ia, on average, to differ from those of low-redshift SNe~Ia. Since all galactic environments at redshift $0 < z < 1.5$ are also present in the local Universe, SN~Ia evolution does not necessarily mean that there are not local analogs. For instance, if the distribution of observables is on average different at high redshift, as long as for each high-redshift SN there is a similar low-redshift counterpart, the peak brightness could, in principle, be correctly translated into an accurate distance. Within the local sample, there is no indication of a correlation between host-galaxy metallicity and light-curve shape \citep{Gallagher05}. In the process of classifying and finding the redshifts for SNe from the ESSENCE (Equation of State: SupErNovae trace Cosmic Expansion) survey \citep{Miknaitis07, Wood-Vasey07}, we have obtained 107 spectra which have accurate light-curve parameters such as $\Delta$ (a light-curve width parameter), time of maximum light, and visual extinction \citep{Matheson05, Foley08a}. Most spectra in this sample have low signal-to-noise ratios (S/N) compared to spectra of low-redshift SNe. This makes impractical a detailed analysis of each object individually to test for outliers. However, by combining the spectra to make composite spectra, we are able to study the mean spectral properties of the samples. In Section~\ref{s:sample} we discuss our low and high-redshift SN~Ia spectral samples. We describe our methods of creating composite spectra in Section~\ref{s:method}. In Section~\ref{s:results} we present the composite spectra and compare the two samples, while in Section~\ref{s:discussion} we discuss the implications of these results. We present our conclusions in Section~\ref{s:conclusions}. Throughout this paper we assume the standard cosmological model with $(h, \Omega_{m}, \Omega_{\Lambda}) = (0.7, 0.3, 0.7)$. | \label{s:conclusions} By combining many low-S/N, high-redshift SN~Ia spectra, we are able to construct the first series of composite SN~Ia spectra based on the parameters of redshift, phase, $\Delta$, and $A_{V}$. In addition, we constructed similar composite spectra from a high-quality sample of low-redshift SN~Ia spectra obtained over the last two decades. Comparison of the composite spectra has shown that once we account for galaxy-light contamination, the two samples are remarkably similar. There are several minor deviations between low and high-redshift samples. These deviations fall into three categories: related to metallicity, related to $^{56}$Ni production, and unknown. The UV excess in the ESSENCE premaximum spectrum is indicative of a different metallicity for the low and high-redshift SNe. Depending on the model, a UV excess is the result of higher or lower metallicity \citep{Hoflich98, Lentz00}. The stronger, more blueshifted \ion{Si}{2} $\lambda 6355$ line in the ESSENCE premaximum spectrum indicates a higher metallicity \citep{Lentz00}. The most significant difference between our samples is the varying strength of the \ion{Fe}{3} $\lambda 5129$ line. The evolution of the \ion{Fe}{3} line may indicate that high-redshift SNe~Ia have lower temperatures than low-redshift SNe~Ia. This, in turn, suggests that SNe~Ia should be less luminous at high redshift. The weak \ion{Fe}{3} line may indicate lower $^{54}$Fe production, which could be the result of lower metallicity. Alternatively, lower metallicity may cause less backwarming from UV photons, decreasing the temperature and the \ion{Fe}{3}/\ion{Fe}{2} ratio. It is unclear if this difference is from evolution in all SNe~Ia, evolution just in the low-$\Delta$ SNe~Ia, changing demographics, or a selection effect. Low-$\Delta$ SNe~Ia tend to come from the short-delay progenitor channel. Because of the short delay, the progenitors of these SNe are more biased by galactic environment, and thus galactic evolution, than the progenitors of long-delay channel SNe~Ia. It is therefore not surprising that the difference in the \ion{Fe}{3} line is more obvious in the low-$\Delta$ objects. It is possible that the different strengths of the \about3000~\AA\ \ion{Fe}{2} feature between low and high redshift is an artifact of the construction of the composite spectra. An analysis of the individual spectra of both samples indicates that the samples are not significantly different; however, the small number of UV objects hampers this study. It is difficult to definitively detect metallicity differences between the Lick and ESSENCE samples. First, we have observed three differences between the samples which harbinger a difference in metallicity: a UV excess, a stronger and more blueshifted \ion{Si}{2} $\lambda 6355$ line, and a weaker \ion{Fe}{3} $\lambda 5129$ line. The UV excess is an ambiguous indicator since the models disagree if it indicates lower or higher metallicity. The \ion{Si}{2} line in the premaximum spectrum suggests higher metallicity, and the \ion{Fe}{3} line suggests lower metallicity. We have also shown that the previously published low-redshift template spectra have multiple drawbacks when comparing to high-redshift composite spectra. We therefore caution against using these templates for studies of SN~Ia evolution. Furthermore, deriving K-corrections from any low-redshift template is difficult; however, the systematic errors are likely to be relatively small. We see that the intrinsic variation of low-redshift SN spectra is \about3\% in the optical. The spectra vary more in the near-UV and UV as suggested by photometry \citep{Jha06}. We are able to put the first constraints of SN~Ia evolution to $\lesssim 10$\%. The results of this study are very suggestive, but require further investigation. In order to improve our understanding of SN~Ia evolution, we propose three future studies related to this work. First, the theoretical models of the effects of metallicity on SN~Ia spectra should be expanded. With the current ambiguity amongst models, we cannot determine the direction of the trend in metallicity. Second, we should gather many more high-redshift spectra to disentangle the redshift-$\Delta$ ambiguity. Finally, further UV observations of nearby SNe~Ia are desperately needed. The only current instrument available for the task is the {\it Swift} \uvot. However, previous attempts at obtaining SN~Ia UV spectra have been disappointing \citep{Brown05}. We suggest an intense campaign spending several hours per spectrum (similar to \iue) with the \uvot. In the near future, we may once again have the ability to obtain high-quality UV spectra with \hst\, using \stis\ or COS. If the upcoming \hst\, servicing mission is successful, we strongly suggest a massive campaign to observe local SNe~Ia in the UV. Since \jwst\, does not have the capabilities to observe the UV, this may be our last opportunity for many years. | 7 | 10 | 0710.2338 |
0710 | 0710.2879_arXiv.txt | % The physical ingredients to describe the epoch of cosmological recombination are amazingly simple and well-understood. This fact allows us to take into account a very large variety of processes, still finding potentially measurable consequences. In this contribution we highlight some of the detailed physics that were recently studied in connection with cosmological hydrogen recombination. The impact of these considerations is two-fold: (i) the associated release of photons during this epoch leads to interesting and {\it unique deviations} of the Cosmic Microwave Background (CMB) energy spectrum {\it from a perfect blackbody}, which, in particular at decimeter wavelength, may become observable in the near future. Observing these distortions, in principle would provide an additional way to determine some of the key parameters of the Universe (e.g. the specific entropy, the CMB monopole temperature and the pre-stellar abundance of helium), {\it not suffering} from limitations set by {\it cosmic variance}. Also it permits us to confront our detailed understanding of the recombination process with {\it direct observational evidence}. In this contribution we illustrate how the theoretical {\it spectral template} for the cosmological recombination spectrum may be utilized for this purpose. (ii) with the advent of high precision CMB data, e.g. as will be available using the {\sc Planck} Surveyor or {\sc Cmbpol}, a very accurate theoretical understanding of the {\it ionization history} of the Universe becomes necessary for the interpretation of the CMB temperature and polarization anisotropies. Here we show that the uncertainty in the ionization history due to several processes that until now are not taken in to account in the standard recombination code {\sc Recfast} exceed the level of $0.1\%$ to $0.5\%$ for each of them. However, it is indeed surprising how {\it inert} the cosmological recombination history is even at percent-level accuracy. | \label{RS:sec:Intro} \subsection{What is so rich and beautiful about cosmological recombination?} \label{RS:sec:Intro1} Within the cosmological concordance model the physical environment during the epoch of cosmological recombination (redshifts $500 \lesssim z\lesssim 2000$ for hydrogen, $1600 \lesssim z\lesssim 3500$ for \ion{He}{II}$\rightarrow$\ion{He}{I} and $5000 \lesssim z\lesssim 8000$ for \ion{He}{III}$\rightarrow$\ion{He}{II} recombination) is extremely simple: the Universe is homogeneous and isotropic, globally neutral and is expanding at a rate that can be computed knowing a small set of cosmological parameters. The baryonic matter component is dominated by hydrogen ($\sim 76\%$) and helium ($\sim 24\%$), with negligibly small traces of other light elements, such as deuterium and lithium, and it is continuously exposed to a bath of isotropic blackbody radiation, which contains roughly $1.6\times 10^9$ photons per baryon. These initially simple and very unique settings in principle allows us to predict the {\it ionization history} of the Universe and the {\it cosmological recombination spectrum} (see Sect.~\ref{RS:sec:spectrum}) with extremely high accuracy, where the limitations are mainly set by our understanding of the {\it atomic processes} and associated transition rates. It is this simplicity that offers us the possibility to enter a {\it rich} field of physical processes and to challenge our understanding of atomic physics, radiative transfer and cosmology, eventually leading to a {\it beautiful} variety of potentially observable effects. \subsection{What is so special about cosmological recombination?} \label{RS:sec:Intro2} The main reason for this simplicity is the {\it extremely large specific entropy} and the {\it slow expansion} of our Universe. Due to the huge number of CMB photons, the free electrons are tightly coupled to the radiation field due to tiny energy exchange during {\it Compton scattering} off thermal electrons until rather low redshifts, such that during recombination the thermodynamic temperature of electrons is equal to the CMB blackbody temperature with very high precision. In addition, the very fast {\it Coulomb interaction} and atom-ion collisions allows to maintain full thermodynamic equilibrium among the electrons, ions and neutral atoms down to $z\sim 150$ \citep{RS_Zeldovich68}. Also, processes in the baryonic sector cannot severely affect any of the radiation properties, down to redshift where the first stars and galaxies appear, and the atomic rates are largely dominated by radiative processes, including {\it stimulated recombination}, {\it induced emission} and absorption of photons. On the other hand, the slow expansion of the Universe allows us to consider the evolution of the atomic species along a sequence of {\it quasi-stationary} stages, where the populations of the levels are nearly in full equilibrium with the radiation field, but only subsequently and very slowly drop out of equilibrium, finally leading to {\it recombination} and the {\it release of additional photons} in uncompensated bound-bound and free-bound transitions. \subsection{Historical overview.} \label{RS:sec:history} It was realized at the end of the 60's \citep{RS_Zeldovich68, RS_Peebles68}, that during the epoch of cosmological hydrogen recombination (typical redshifts $800\lesssim z \lesssim 1600$) any direct recombination of electrons to the ground state of hydrogen is immediately followed by the ionization of a neighboring neutral atom due to re-absorption of the newly released Lyman-continuum photon. In addition, because of the enormous difference in the $2{\rm p}\leftrightarrow 1{\rm s}$ dipole transition rate and the Hubble expansion {rate}, photons emitted close to the center of the Lyman-$\alpha$ line scatter $\sim 10^8-10^9$ times before they can finally escape further interaction with the medium and thereby permit a successful settling of electrons in the 1s-level. It is due to these very peculiar circumstances that the $2{\rm s}\leftrightarrow 1{\rm s}$-two-photon decay process, being $\sim 10^8$ orders of magnitude slower than the Lyman-$\alpha$ resonance transition, is able to substantially control the dynamics of cosmological hydrogen recombination \citep{RS_Zeldovich68, RS_Peebles68}, allowing about 57\% of all hydrogen atoms in the Universe to recombine at redshift $z\lesssim 1400$ through this channel \citep{RS_Chluba2006b}. Shortly afterwards \citep{RS_Sunyaev1970, RS_Peebles1970}, the importance of the ionization history as one of the key ingredients for the theoretical predictions of the Cosmic Microwave Background (CMB) temperature and polarization anisotropies became clear, and today these tiny directional variations of the CMB temperature ($\Delta T/T_0\sim 10^{-5}$) around the mean value $T_0=2.725\pm 0.001\,$K \citep{RS_Fixsen2002} have been observed for the whole sky using the {\sc Cobe} % and {\sc Wmap} % satellites, beyond doubt with great success. The high quality data coming from balloon-borne and ground-based CMB experiments ({\sc Boomerang, Maxima, Archeops, Cbi, Dasi} and {\sc Vsa} etc.) today certainly provides one of the mayor pillars for the {\it cosmological concordance model} \citep{RS_Bennett2003}, and planned CMB mission like the {\sc Planck} Surveyor will help to further establish the {\it era of precision cosmology}. \begin{figure} \centering \includegraphics[width=0.8\columnwidth]{./RS.plot.1.eps} \caption{Ionization history of the Universe and the origin of CMB signals. The observed angular fluctuations in the CMB temperature are created close to the maximum of the Thomson visibility function around $z\sim 1089$, whereas the direct information carried by the photons in the cosmological hydrogen recombination spectrum is from earlier times.} \label{RS:fig:plot1} \end{figure} In September 1966, one of the authors (RS) was explaining during a seminar at the Shternberg Institute in Moscow how according to the Saha formula this recombination should occur. After the talk his friend (UV astronomer) {\it Vladimir Kurt} asked him: {\it 'but where are all the redshifted Lyman-$\alpha$ photons that were released during recombination?'} Indeed this was a great question, which was then addressed in detail by \citet{RS_Zeldovich68}, leading to an understanding of the role of the 2s-two-photon decay, the delay of recombination as compared to the Saha-solution, the spectral distortions of the CMB due to two-photon continuum and Lyman-$\alpha$ emission, the frozen remnant of ionized atoms, and the radiation and matter temperature equality until $z\sim 150$. All recombined electrons in hydrogen lead to the release of $\sim 13.6\,$eV in form of photons, but due to the large specific entropy of the Universe this will only add some fraction of $\Delta \rho_\gamma/\rho_\gamma\sim 10^{-9}-10^{-8}$ to the total energy density of the CMB spectrum, and hence the corresponding distortions are expected to be very small. However, all the photons connected with the Lyman-$\alpha$ transition and the 2s-two-photon continuum today appear in the Wien part of the CMB spectrum, where the number of photons in the CMB blackbody is dropping exponentially, and, as realized earlier \citep{RS_Zeldovich68, RS_Peebles68}, these distortions are significant (see Sect.~\ref{RS:sec:spectrum}). In 1975, {\it Victor Dubrovich} pointed out that the transitions among highly excited levels in hydrogen are producing additional photons, which after redshifting are reaching us in the cm-spectral band. This band is actually accessible from the ground. Later these early estimates were significantly refined by several groups (e.g. see \citet{RS_Kholu2005} and \citet{RS_Jose2006} for references), with the most recent calculation performed by \citet{RS_Chluba2006b}, also including the previously neglected free-bound component, and showing in detail that the relative distortions are becoming more significant in the decimeter Rayleigh-Jeans part of the CMB blackbody spectrum (Fig.~\ref{RS:fig:DI_results}). These kind of precise computations are becoming feasible today, because (i) our knowledge of atomic data (in particular for neutral helium) has significantly improved, and (ii) it is now possible to handle large systems of strongly coupled differential equations using modern computers. The most interesting aspect of this radiation is that it has a very {\it peculiar} but {\it well-defined, quasi-periodic spectral dependence}, where the photons are coming from redshifts $z\sim 1300-1400$, i.e. {\it before} the time of the formation of the CMB angular fluctuation close to the maximum of the Thomson visibility function (see Fig.~\ref{RS:fig:plot1}). Therefore, measuring these distortions of the CMB spectrum would provide a way to confront our understanding of the recombination epoch with {\it direct experimental evidence}, and in principle may open another independent way to determine some of the key parameters of the Universe, in particular the value of the CMB monopole temperature, $T_0$, the number density of baryons, $\propto \Omega_{\rm b}h^2$, or alternatively the specific entropy, and the primordial helium abundance (e.g. see \citet{RS_Chluba2007d} and references therein). | \label{RS:sec:conclusion} It took several decades until measurements of the CMB temperature fluctuations became a reality. After {\sc Cobe} the progress in experimental technology has accelerated by orders of magnitude. Today CMB scientists are even able to measure $E$-mode polarization, and the future will likely allow to access the $B$-mode component of the CMB in addition. Similarly, one may hope that the development of new technologies will render the consequences of the discussed physical processes observable. Therefore, also the photons emerging during the epoch of cosmological (hydrogen) recombination could open another way to refine our understanding of the Universe. As we illustrated in this contribution, by observing the CMB spectral distortions from the epoch of cosmological recombination we can in principle directly measure cosmological parameters like the value of the CMB monopole temperature, the specific entropy, and the pre-stellar helium abundance, {\it not suffering} from limitations set by {\it cosmic variance}. Furthermore, we could directly test our detailed understanding of the recombination process using {\it observational data}. It is also remarkable that the discussed CMB signal is coming from redshifts $z\sim 1300-1400$, i.e. before bulk of the CMB angular fluctuations were actually formed. To achieve this task, {\it no absolute measurement} is necessary, but one only has to look for a modulated signal at the $\sim\mu$K level, with typical amplitude of $\sim 30\,$nK and $\Delta \nu/\nu\sim 0.1$, where this signal can be predicted with high accuracy, yielding a {\it spectral template} for the full cosmological recombination spectrum, also including the contributions from helium. | 7 | 10 | 0710.2879 |
0710 | 0710.0798_arXiv.txt | We demonstrate that low resolution Ca\,{\small II} triplet (CaT) spectroscopic estimates of the overall metallicity ([Fe/H]) of individual Red Giant Branch (RGB) stars in two nearby dwarf spheroidal galaxies (dSphs) agree to $\pm$0.1-0.2 dex with detailed high resolution spectroscopic determinations for the same stars over the range $-2.5 <$ [Fe/H] $< -0.5$. For this study we used a sample of 129 stars observed in low and high resolution mode with VLT/FLAMES in the Sculptor and Fornax dSphs. We also present the data reduction steps we used in our low resolution analysis and show that the typical accuracy of our velocity and CaT [Fe/H] measurement is $\sim$2 \kms and 0.1 dex respectively. We conclude that CaT-[Fe/H] relations calibrated on globular clusters can be applied with confidence to RGB stars in composite stellar populations over the range $-2.5 <$ [Fe/H] $< -0.5$. | An important aspect for a full understanding of galactic evolution is the metallicity distribution function of the stellar population with time. Carrying out detailed abundance analyses with high resolution (HR) spectroscopy to trace the patterns that allow one to distinguish between the different galactic chemical enrichment processes is time consuming for large samples of individual stars in a galaxy. This is partly due to the observing time required, but also because of the complex data reduction and analysis necessary. Fortunately, there is an empirically developed, simply calibrated method available which can make an efficient estimate of metallicity ([Fe/H]) for individual Red Giant Branch (RGB) stars using the strength of the Ca\,{\small II} triplet (CaT) lines at 8498, 8542, 8662 \AA. This method was pioneered for use on individual stars by \citet{arm1991}. It has the advantage that the lines are broad enough that they can be accurately measured with moderate spectral resolution \citep[e.g,][]{cole2004}. The CaT method is routinely used to estimate [Fe/H] for nearby resolved stellar systems and also provides an accurate radial velocity estimate. Both measurements are facilitated by the strength of the CaT lines and by the generally red colours of the target stars. However, the CaT-derived abundances are empirically defined, with a poorly understood physical basis. Therefore it is important to check the results against HR spectroscopic (\ie direct) measurements of [Fe/H] and other elements. The ``classical'' CaT calibration is based on the use of globular cluster stars, all which are drawn from a single age and metallicity stellar population. The CaT equivalent widths are directly compared to HR spectroscopic measurements of [Fe/H] over a range of metallicity, and this comparison is used to define the relation between CaT equivalent width (EW) and [Fe/H] for all observations taken with the same set up. This approach has been extensively tested for a large sample of globular clusters \citep[and references therein]{rutledge1997a, rutledge1997b}. However, globular clusters typically exhibit a constant [Ca/Fe] for a large range of [Fe/H]. This leads to uncertainty in the effect of varying [Ca/Fe] ratios such as is seen in the more complex stellar populations found in galaxies. Furthermore, stars in dwarf galaxies invariably cover a significant range of ages as well as metallicities. This mismatch in the properties of calibrators and targets has led to suggestions that the CaT method may not be a very accurate indicator of [Fe/H] for more complex stellar populations, especially in those cases where [Ca/Fe] varies significantly \citep[e.g,][]{pont2004}. In this paper we investigate the validity of the CaT method for complex stellar populations. We compare large samples of [Fe/H] measurements coming from VLT/FLAMES made using both the CaT method and direct HR spectroscopic measurements {\it for the same stars} in two nearby dwarf spheroidal galaxies (dSphs), Sculptor and Fornax, over a range of [Fe/H] and [Ca/H]. This is the first time such a detailed comparison has been made for stars outside globular clusters. We also investigate the theoretically predicted behaviour of the CaT method for a range of stellar atmospheric parameters using a grid of model atmosphere spectra from \citet{munari2005}. The paper is organised as follows. In Section~\ref{sect:datareduction} we describe the data reduction steps we use within the DART (Dwarf galaxy Abundances and Radial velocities Team) collaboration to estimate EW and velocities from observations in the CaT region, as the accuracy with which this can be done clearly has important implications for the reliability of our conclusions for these galaxies. We also discuss the verification of the overall calibration and accuracy of the velocity and EW measurements by comparison of results from independent CaT observations and by comparing with theoretical expectations based on signal-to-noise, resolution and line profile properties. In Section~\ref{sec:calibration_gcs} we derive the standard CaT-[Fe/H] globular cluster calibration for low resolution (LR) VLT/FLAMES data. In Section~\ref{sec:calibration_hr} we compare the derived [Fe/H] from the CaT to the HR [Fe/H] for the Sculptor and Fornax dSphs. Finally, in Section~\ref{sec:uncert} we discuss the uncertainties that come from using Ca\,{\small II} lines to derive an [Fe/H] abundance for stellar populations where the $\alpha$-abundance varies and use a comparison with stellar model atmospheres to further investigate age, metallicity and $\alpha$-abundance effects. | We described the data reduction steps we use within the DART collaboration to estimate velocities and CaT metallicities for LR data for RGB stars in dSphs. We showed that we obtain accurate velocities and [Fe/H] measurements, with internal error in velocity $\sim$ 2 \kms and in [Fe/H] $\sim$ 0.1 dex at a S/N per \AA\, of 20. We used 4 Galactic globular clusters observed with VLT/FLAMES in the CaT region to test the performance of several CaT $W'$-[Fe/H] relations existing in the literature. We also derived the best calibration from these globular cluster data. The relation here derived is consistent at the 1-$\sigma$ level with the calibration derived in \citet{tolstoy2001}, which we used in \citet{tolstoy2004}, \citet{battaglia2006} and \citet{helmi2006}. We used a sample of 93 and 36 RGB stars in the Sculptor and Fornax dSphs, respectively, overlapping between our LR and HR VLT/FLAMES observations, to test the globular clusters CaT calibration. This is the first time that the CaT calibration is tested on field stars in galaxies. We find a good agreement between the metallicities derived with these two methods. However, a systematic trend is present with [Fe/H], such that using the globular cluster calibration derived in this work the [Fe/H] measurement from CaT is overestimated of $\sim$0.1 dex at [Fe/H]$_{\rm HR} < -2.2$, whilst at [Fe/H]$_{\rm HR} > -1.2$ it is underestimated of $\sim$0.1-0.2 dex. No clear systematic trend is instead derived from our data for [Fe/H]$_{\rm HR} > -0.8$. In order to understand this systematic effect, we explored the possible contribution of Ca abundance to the calibration, and showed that there are indications that it might well affect the CaT $W'$, although much less than [Fe/H]. From our dataset we also show that, contrary to previous claims, it is not advisable to use the CaT $W'$ as a linear indicator of [Ca/H]. Finally we investigated the effect of varying stellar atmosphere parameters on the CaT method by analysing a large sample of model atmosphere spectra \citep{munari2005}. We again demonstrated that the CaT method is (surprisingly) robust to the usual combination of age, metallicities and [$\alpha$/Fe] variations seen in nearby dSphs. From our analysis we see that even for large differences in [Ca/Fe] between calibrating globular clusters and our sample of dSphs ($\sim$ 0.5 dex at [Fe/H]$\sim -0.7$) and large difference in ages (at [Fe/H] $\sim -0.7$ Fornax stars are $\sim$ 10 Gyr younger than globular clusters stars) the error in estimating [Fe/H] using globular clusters as calibrators is just 0.1-0.2 dex. The Fornax dSph is likely to represent the most extreme case as it has had one of the most extended star formation histories among the dSphs in the Local Group. We conclude that CaT-[Fe/H] relations calibrated on globular clusters can be applied with confidence to RGB stars in composite stellar populations such as galaxies, at least in the [Fe/H] range probed by the above analyses, $-2.5 <$ [Fe/H] $<-0.5$. Hence, the CaT method provides a good indicator of the overall metallicity of resolved stars. This has implications for the efficiency with which we can obtain metallicity distribution functions of nearby resolved galaxies. The HR data collected in this paper required more than 6 nights of VLT observing time for $\sim$ 150 spectra, whereas $\sim 120$ CaT spectra were obtained in one hour VLT observing time. | 7 | 10 | 0710.0798 |
0710 | 0710.1082_arXiv.txt | We report a new determination of the faint end of the galaxy luminosity function in the nearby clusters Virgo and Abell 2199 using data from SDSS and the Hectospec multifiber spectrograph on the MMT. The luminosity function of A2199 is consistent with a single Schechter function to $M_r=-15.6$ + 5 log $h_{70}$ with a faint-end slope of $\alpha=-1.13\pm0.07$ (statistical). The LF in Virgo extends to $M_r\approx-13.5\approx M^*+8$ and has a slope of $\alpha=-1.28\pm0.06$ (statistical). The red sequence of cluster members is prominent in both clusters, and almost no cluster galaxies are redder than this sequence. A large fraction of photometric red-sequence galaxies lie behind the cluster. We compare our results to previous estimates and find poor agreement with estimates based on statistical background subtraction but good agreement with estimates based on photometric membership classifications (e.g., colors, morphology, surface brightness). We conclude that spectroscopic data are critical for estimating the faint end of the luminosity function in clusters. The faint-end slope we find is consistent with values found for field galaxies, weakening any argument for environmental evolution in the relative abundance of dwarf galaxies. However, dwarf galaxies in clusters are significantly redder than field galaxies of similar luminosity or mass, indicating that star formation processes in dwarfs do depend on environment. | The luminosity function of galaxies is fundamental to understanding galaxy formation and evolution. The luminosity function differs dramatically from the expected mass function of dark matter halos, indicating that baryonic physics is very important for understanding galaxies. In particular, a well-determined luminosity function enables accurate modeling linking the masses of dark matter haloes to galaxy luminosities \citep[e.g.,][and references therein]{vale06,yang07}. These empirical models provide a powerful test of any model of galaxy formation and evolution. Early studies of the luminosity function used the large galaxy density in clusters as a tool for measuring the shape of the luminosity function \citep[e.g.,][]{sandage85}. The obvious drawback of this method is that the luminosity function in dense environments may differ from that in more typical galaxy environments \citep[][]{binggeli88,driver94,depropris95}. Environmental trends in the luminosity function may reflect differences in galaxy formation in different environments \citep[][]{tully02,benson03}. For instance, tidal stripping or ``threshing'' of larger galaxies may produce dwarf galaxies \citep{bekki01}, or dwarf galaxies may be formed in tidal tails of intractions among giant galaxies \citep{barnes92}. Alternatively, the denser environments of protoclusters may have shielded low-mass galaxies from the ultraviolet radiation responsible for reionization \citep{tully02,benson03}. Many studies suggest an environmental influence on the LF; others provide no such evidence. The main difficulty in resolving this important issue is the challenge of determining cluter membership for faint galaxies where background galaxy counts are large. Because few deep spectroscopic surveys of clusters extend into the dwarf galaxy regime \citep[$M_r\gtrsim$-18; for exceptions, see][and references therein]{mobasher03,christlein03,mahdavi05}, cluster membership is usually determined via statistical subtraction of background galaxies \citep[e.g.,][and references therein]{popesso06b,jenkins07,milne07,adami07,yamanoi07,barkhouse07}. Because galaxy number counts increase much more steeply than cluster member counts (even for very steep faint-end slopes), small systematic uncertainties in background subtraction can produce large uncertainties in the abundance of faint cluster galaxies. Here, we use MMT/Hectospec spectroscopy and data from the Sloan Digital Sky Survey \citep[SDSS,][]{sdss} to estimate the luminosity function (LF) in the clusters Abell 2199 and Virgo. These data enable very deep sampling of the luminosity function. In particular, we report an estimate of the faint-end slope of the luminosity function with much smaller systematic uncertainties than most previous investigations. We demonstrate that photometric properties of galaxies such as color and surface brightness correlate well with cluster membership (in agreement with many previous studies). Very few galaxies redder than the red sequence are cluster members. We discuss the photometric and spectroscopic data in $\S 2$. We present the luminosity functions in $\S 3$. We compare our results to previous studies and discuss possible systematic effects and uncertainties in $\S 4$. We conclude in $\S 5$. An Appendix details the construction of our catalog of confirmed and probable Virgo cluster members. We assume cosmological parameters of $\Omega_m$=0.3, $\Omega_\Lambda$=0.7, $H_0$=70~$h_{70} \kms \mbox{Mpc}^{-1}$. The spatial scale at the distance of A2199 is 1\arcs=0.61~$\kpc$ and 1\arcs=0.080~$\kpc$ at the distance of Virgo. | Determining the faint end of the luminosity function in clusters has remained an unresolved problem for many years. We report a new estimate of the luminosity function in A2199 from MMT/Hectospec spectroscopy, and of the Virgo cluster from SDSS DR6 data. Both LFs extends to fainter absolute magnitudes than most previous determinations of the LF in massive clusters. The LF closely follows a Schechter function to $M_r\approx-15\approx M^*+6$ in A2199 and to $M_r\approx-13\approx M^*+8$ in Virgo. There are no large populations of high surface brightness galaxies or galaxies redder than the red sequence that contribute significantly to the LF in either cluster. In A2199, we find that essentially no galaxies redward of the photometric red sequence are cluster members. However, the red sequence itself contains a significant number of background galaxies. We find no evidence of an upturn in the A2199 LF at faint magnitudes as claimed by some recent studies \citep[][]{popesso06b,jenkins07,milne07,adami07,yamanoi07,barkhouse07}, although the range of absolute magnitudes precludes a conclusive result. A simple extrapolation of the A2199 LF using the membership fractions at the spectroscopic limit provides an approximate upper limit of the LF at fainter magnitudes; this extrapolation suggests that the faint end slope is probably comparable to that of Virgo. In the Virgo cluster, we find an LF consistent with a moderate faint-end slope ($\alpha=-1.28\pm0.06$). The Virgo LF extends much fainter than the A2199 LF, and we conclusively demonstrate that the Virgo LF is inconsistent with the steeper LFs found by \citet{popesso06b}. The discrepancy between the LFs may be due to systematic uncertainties in statistical background subtraction, and we discuss some possibilities. Other estimates of the LF using deep spectroscopy find slopes similar to ours \citep{mobasher03,christlein03,mahdavi05}. Recent estimates of the LF in Fornax using surface brightness as a membership classification \citep{hilker03} or surface brightness fluctuations to estimate distances \citep{mieske07} indicate a shallow LF similar to the Virgo SDSS LF. Similarly, a clever application of gravitational lensing in A1689 by \citet{medezinski07} suggests a LF consistent with those we find in Virgo and A2199. Low surface brightness galaxies remain problematic for determining cluster LFs. Their low surface brightnesses prohibit reliable redshift estimates using SDSS DR6 spectroscopy. Perhaps the most significant impact of the SDSS spectra is to demonstrate conclusively that higher surface brightness galaxies are virtually all in the background of Virgo. A simple division in apparent magnitude versus surface brightness is a surprisingly powerful membership classification (Figure \ref{virgolsb}, $\S 2.3$). Careful inspection of SDSS galaxies failing the spectroscopic target selection criteria reveals many low surface brightness galaxies that are likely Virgo members. It is somewhat ironic that almost all $r<17.5$ galaxies within 1 Mpc of M87 without well-measured redshifts are the LSB galaxies most likely to be Virgo members. We list these galaxies in the Appendix as an aid for future efforts to obtain spectroscopy of these galaxies. The spectroscopically determined LFs we find for the A2199 and Virgo clusters are similar to the field LF, an important result for models of galaxy formation. Future studies of the LF in more clusters and at larger clustrocentric radii will constrain cluster-to-cluster variations in the LF and any radial dependence. The fraction of galaxies belonging to a cluster decreases dramatically with both increasing magnitude and increasing projected radius. Even at the relatively small radii we probe here, we show that statistical background subtraction is problematic due to the high precision required. We therefore recommend that future investigations of the LF in clusters avoid statistical background subtraction and instead identify member galaxies via spectroscopy or photometric information such as colors or surface brightness fluctuations. | 7 | 10 | 0710.1082 |
0710 | 0710.3614_arXiv.txt | Ultraviolet, optical and near infrared images of the nearby ($D \approx 5.5$ Mpc) SBm galaxy NGC 1311, obtained with the {\it Hubble Space Telescope}, reveal a small population of 13 candidate star clusters. We identify candidate star clusters based on a combination of their luminosity, extent and spectral energy distribution. The masses of the cluster candidates range from $\sim$10$^3$ up to $\sim$10$^5$ $M_{\odot}$, and show a strong positive trend of larger mass with increasing with cluster age. Such a trend follows from the fading and dissolution of old, low-mass clusters, and the lack of any young super star-clusters of the sort often formed in strong starbursts. The cluster age distribution is consistent with a bursting mode of cluster formation, with active episodes of age $\sim$10 Myr, $\sim$100 Myr and $\ga$1 Gyr. The ranges of age and mass we probe are consistent with those of the star clusters found in quiescent Local Group dwarf galaxies. | Star clusters are powerful tools for probing the star-formation history and chemical enrichment of galaxies (e.g., \citeauthor{h61}\citeyear{h61}; \citeauthor{whit99}\citeyear{whit99}; \citeauthor{dk2}\citeyear{dk2}; \citeauthor{grijs}\citeyear{grijs}). In nearby systems, star clusters provide us with spatially resolved examples of simple stellar populations (SSPs). Observations of individual cluster stars are, therefore, crucial for understanding the effect of metallicity on stellar evolution (e.g., \citeauthor{carme}\citeyear{carme} and references therein). In more distant systems, star clusters provide us with examples of unresolved SSPs. Analysis of statistical samples of such clusters offers us a wealth of information on the history of star formation in their parent galaxies (e.g., \citeauthor{grijs}\citeyear{grijs} and references therein). Understanding star and cluster formation in a range of environments is essential for understanding galaxy evolution because star cluster formation traces the strongest episodes of galactic star formation (e.g., \citeauthor{dGNA}\citeyear{dGNA}). There has been a great deal of work on star cluster populations and formation histories in nearby luminous spiral galaxies over the last decade (e.g., \citeauthor{whit99}\citeyear{whit99}; \citeauthor{grijs}\citeyear{grijs}; \citeauthor{dk2}\citeyear{dk2}; \citeauthor{lar04}\citeyear{lar04}). By comparison, the star cluster formation history in nearby, low-luminosity galaxies has received relatively little attention (but see, e.g., \citeauthor{bhe}\citeyear{bhe}; \citeauthor{a04}\citeyear{a04}; \citeauthor{seth04}\citeyear{seth04}; \citeauthor{dGA06}\citeyear{dGA06}). Studies of nearby galaxies are important because we can resolve abundant small-scale detail in them. Moreover, the low metallicities of low-luminosity nearby galaxies offer us a close-up view of the star-formation process that may better resemble that in the high-redshift, early Universe. Studies of detailed stellar formation histories are still restricted to fairly nearby galaxies, and are most powerful when they combine high angular resolution with broad wavelength coverage. The number of very nearby galaxies with such data available is still quite small. NGC 1311 is a very nearby ($D \approx 5.5$ Mpc), but little-studied late-type (SBm) galaxy. \citeauthor{tully}(\citeyear{tully}) identify it as a member of the 14$+$14 Association, a loose group dominated by the luminous spiral NGC 1313. Table 1 summarizes the basic properties of the system. NGC 1311 was a target in several {\it Hubble Space Telescope} (HST) snapshot surveys (GO programs 9124 and 9824; \citeauthor{wind02}\citeyear{wind02}, \citeauthor{vio}\citeyear{vio}, and \S 2 below). As a result, a set of broad-band images spanning a wide wavelength interval (0.3---1.6$\mu$m) at sub-arcsecond resolution now exists for this galaxy. As NGC 1311 is quite nearby, the bright star clusters and luminous individual stars are detected as discrete sources. We can thus probe the spatially resolved star-formation history of NGC 1311 by studying the broad-band spectral energy distributions of the star clusters, the individual stars, and the unresolved light. This paper is concerned with the star clusters of NGC 1311. We shall address the individual stars and unresolved light in future publications. In \S 2 we describe the HST observational data. We present the observed properties of the candidate star clusters in \S 3, and our analysis of these observations in \S 4. We summarize our conclusions, and discuss issues for further research in \S 5. | NGC 1311 is a nearby, but little studied, low-luminosity late-type spiral. It has optical and \hi\ properties (see Table 1) typical for such galaxies (\citeauthor{m98}\citeyear{m98}; \citeauthor{lvz}\citeyear{lvz}). It is a member of a loose association of galaxies (\citeauthor{tully}\citeyear{tully}), but displays no obvious sign of any recent interaction. We have used HST WFPC2 and NICMOS images of NGC 1311, that span the near-UV through the near-IR, to identify a small population of 13 candidate star clusters. Their masses increase systematically with cluster age, as would follow from the fading and dissolution of old, low-mass clusters, and the lack of any young super star clusters associated with strong starbursts. Half the cluster candidates are significantly fainter than the turnover of the globular cluster luminosity function. We are thus probing the range of luminosities typical of the faint star clusters found in Local Group dwarf galaxies, and the open cluster population of the Galactic disk. Analysis of the photometry of the candidate star clusters suggests that NGC 1311 has had three cluster-forming episodes in its history, occuring $\sim$10 Myr, $\sim$100 Myr, and $\ga$1 Gyr ago. This is consistent with observational work on other nearby low-luminosity galaxies indicating a bursting mode of star-formation. The recent star formation, as traced by the NUV continuum (see Fig.~1a), is concentrated at the east and west ends of the central bulge-like concentration. This is reminiscent of stochastic star formation models (\citeauthor{ssg}\citeyear{ssg}) as well as the observed properties of other low-luminosity star-forming galaxies (e.g., Sextans A; \citeauthor{rdp}\citeyear{rdp}). NGC 1311 is an excellent example of a nearby, low-luminosity, star-forming, gas-rich galaxy that is evolving in relative isolation. Understanding the star- and cluster-formation history and chemical evolution of such galaxies is an essential part of unraveling the problem of galaxy evolution. Our next step is a study of the resolved stellar populations of NGC 1311 with our combined WFPC2/NICMOS HST data. The large wavelength range of our data allows us to sample both the very recent star formation (dominating the UV light) and the ancient stellar populations (from the red/near-IR light) that appear ubiquitous even in very late-type galaxies (e.g., \citeauthor{baa}\citeyear{baa}; \citeauthor{vio}\citeyear{vio}). This should give us a clearer picture of the star formation history in this system, and a fuller understanding of the process of star formation in low-luminosity late-type galaxies in general. | 7 | 10 | 0710.3614 |
0710 | 0710.4202_arXiv.txt | We have carried out sub-mm $^{12}$CO($J=3-2$) observations of 6 giant molecular clouds (GMCs) in the Large Magellanic Cloud (LMC) with the ASTE 10m sub-mm telescope at a spatial resolution of 5 pc and very high sensitivity. We have identified 32 molecular clumps in the GMCs and revealed significant details of the warm and dense molecular gas with $n$(H$_2$) $\sim$ 10$^{3-5}$ cm$^{-3}$ and $T_{\mathrm{kin}}$ $\sim$ 60 K. These data are combined with $^{12}$CO($J=1-0$) and $^{13}$CO($J=1-0$) results and compared with LVG calculations. The results indicate that clumps we detected are distributed continuously from cool ($\sim$ 10 -- 30 K) to warm ($\sim$ higher than 30 -- 200 K), and warm clumps are distributed from less dense ($\sim$ 10$^3$ cm$^{-3}$) to dense ($\sim$ 10$^{3.5 - 5}$ cm$^{-3}$).% We found that the ratio of $^{12}$CO($J=3-2$) to $^{12}$CO($J=1-0$) emission is sensitive to and is well correlated with the local \Halpha flux. We interpret that differences of clump propeties represent an evolutionary sequence of GMCs in terms of density increase leading to star formation.Type I and II GMCs (starless GMCs and GMCs with HII regions only, respectively) are at the young phase of star formation where density does not yet become high enough to show active star formation and Type III GMCs (GMCs with HII regions and young star clusters) represents the later phase where the average density is increased and the GMCs are forming massive stars. The high kinetic temperature correlated with \Halpha flux suggests that FUV heating is dominant in the molecular gas of the LMC. | It is of a fundamental importance in astronomy to understand the evolution of galaxies. Since a major constituent of galaxies is stars, the formation of stars is a fundamental process in galactic evolution. The properties of stars characterize the basic contents of galaxies and their time evolution. We understand from studies of the Milky Way galaxy that giant molecular clouds (GMCs), whose mass ranges from 10$^5$ to 10$^7$ \Msun, are the principal sites of star formation and that it perhaps holds the true in other galaxies as well. We also recognize that the GMC properties (e.g., $L_{\mathrm{CO}}$ -- line width relation, index of mass spectrum) are similar among five galaxies in the Local Group according to the spatially resolved studies \citep{Blitz2006}. This supports the idea that studies of GMCs will be useful in understanding the fundamentals of galactic evolution through the formation and evolution of GMCs and star formation therein. Observational studies of GMCs have been most effectively made by the mm interstellar carbon monoxide emission line at 2.6 mm which allows us to probe molecular gas whose density is greater than $\sim$ 100 cm$^{-3}$. We note that the most abundant species, molecular hydrogen, does not have appropriate line emissions in the mm and sub-mm region due to its zero permanent electric dipole moment and large separation between the lowest energy levels, which are not excited significantly in the typical physical conditions of molecular clouds. Recent advances in sub-mm observations have allowed us to determine physical parameters of molecular clouds over much larger ranges than in the mm region by comparing line intensities between different transitions. These sub-mm studies were initiated by the SEST 15m telescope in Chile followed by instruments in Hawaii, Mauna Kea and in the Swiss Alps at an altitude range from 3700 m to 4200 m, including the CSO 10m, JCMT 15m, and KOSMA 3m telescopes, and the AST/RO 1.6m telescope in Antarctica. Subsequently, in the 2000's, the developements of new instruments at an altitude of $\sim$ 5000 m in Atacama in northern Chile resulted in a superior capability because of the high altitude and dry characteristics of the site. The instruments installed in Atacama include the ASTE 10m, APEX 12m and NANTEN2 4m telescopes. All these instruments are beginning to take new molecular data with significantly better quality than before in terms of noise level as well as angular resolution. It is also noteworthy that the current frequency coverage extends as high as the 800GHz band and even the THz region. Among nearby galaxies we can observe at reasonably high spatial resolutions, the Large and Small Magellanic Clouds offer us a unique opportunity to achieve the highest resolutions due to their unrivaled closeness, 50 -- 60 kpc. In particular, the Large Magellanic Cloud (LMC) is actively forming stars in clusters and is an ideal laboratory for us to study star formation, particularly massive star formation in star cluster. In the LMC, the metallicity is a factor of $\sim$ 3 lower than in the Solar neighborhood \citep{Dufour1982, Dufour1984, Rolleston2002}. Also, the visual extinctions are lower and the FUV field is stronger in the LMC than in the Milky Way \citep{Israel1986}, characterizing the initial conditions of star formation. The first spatially resolved complete sample of GMCs in a single galaxy has been obtained towards the whole LMC with the NANTEN 4m telescope in 2.6 mm CO emission at 40 pc resolution \citep{Fukui1999, Fukui2001, Fukui2006b, Mizuno2001}. These studies revealed the three types of GMCs in terms of star formation activities; Type I is starless, Type II is with HII regions only, and Type III is associated with active star formation indicated by huge HII regions and young star clusters, where the stars identified are only O stars due to the sensitivity limitation. It also revealed that the lifetime of a GMC is as short as $\sim$ 30 Myrs \citep{Fukui2006a, Kawamura2006, Kawamura2007}. These previous studies naturally place the LMC as one of the prime targets for sub-mm studies to derive the physical parameters of GMCs. Another aspect which deserves our attention is that very young, rich stellar clusters are forming in the LMC. These are so called populous clusters which are very rare in the Milky Way and resemble globulars formed in the primeval Milky Way. The open clusters forming in the Milky Way are small in the number of stars and loose in spatial distribution. Along with the low metallicity of the LMC, it is an interesting possibility to use molecular data to investigate the formation mechanism of super star clusters at the molecular cloud stage. In the past, there have been some studies that used the higher transitions ($J=2-1$, $J=3-2$, $J=4-3$, $J=7-6$) of CO spectra of the molecular clouds in the LMC \citep[e.g.,][]{Sorai2001,Johansson1998,Heikkila1999,Bolatto2005,Israel2003,Kim2004,Kim2006}. These studies suggest the molecular gas may be warmer and/or denser than in the Milky Way. \citet{Johansson1998} used the SEST 15m telescope to observe the central part of the 30 Doradus nebula (rms $\sim$ 0.2 K in 0.5 km s$^{-1}$ velocity resolution for $J=1-0$ and rms $\sim$ 1.0 K in 0.5 km s$^{-1}$ velocity resolution for $J=3-2$), and the southern HII regions N 158C, N 159, and N 160 with a few prominent CO clouds in the $J=2-1$ and $J=3-2$ transitions of CO. They find that the kinetic temperatures are 10 -- 80 K and the highest temperature is towards 30 Dor. The smallest beam size and grid spacing are 15$\arcsec$ and 11$\arcsec$ respectively in the $J=3-2$ emission. \citet{Heikkila1999} used SEST to observe the $J=3-2$ transition of CO in N 159 and 30 Doradus, as well as other rarer molecular species. This study aimed at obtaining chemical abundance, while it also provides more information on cloud temperature etc, from CO($J=3-2$) data. The kinetic temperatures they derived are 50 K in 30 Dor-10, 15 K in 30 Dor-27, and 20 -- 25 K in N 159W and N 160. \citet{Bolatto2005} employed the AST/RO to observe the $^{12}$CO($J=4-3$) transition at 461 GHz with a 109$\arcsec$ beam. They observed 9 regions in the LMC at 6$\arcmin$$\times$6$\arcmin$ field all with HII regions and derived kinetic temperatures from a comparison between the CO($J=4-3$) and ($J=1-0$) transitions. N 48, N 55A, N 79, N 83A, N 113, N 159W, N 167, N 214C, and LIRL 648 are included. They derive temperatures of 100 -- 300 K and note a trend that higher temperatures occur in moderate density regions, 100 -- 1000 cm$^{-3}$, and the lower temperatures in much denser regions 10$^{4-5}$ cm$^{-3}$. These studies were preceded by a suggestion that significant amounts of warm molecular gas may exist in the LMC \citep{Israel2003}. \citet{Kim2004} also made similar observations towards an HII region, N 44, and suggest very dense gas of $\sim$ 10$^5$ cm$^{-3}$. Most recently, \citet{Kim2006} derived $T_{\mathrm{kin}}$ = 100 K, $n$ $\sim$ 10$^{4.3}$ cm$^{-3}$ for 30 Dor from the intensity ratios of $^{12}$CO($J=7-6$) to $^{12}$CO($J=4-3$) and $^{12}$CO($J=1-0$) to $^{13}$CO($J=1-0$). In the present study, we aim to obtain sub-mm molecular data at better S/N ratios than in the previous studies to make estimates of temperatures and densities over a large sample in the LMC. We will combine the $^{12}$CO($J=3-2$) data obtained with the ASTE telescope and CO($J=1-0$) data obtained with the SEST and Mopra telescopes. In order to make reasonable comparisons between the two transitions, $J=3-2$ and $J=1-0$, we shall smooth the ASTE results (22$\arcsec$ beam) to the same resolution as the SEST data (45$\arcsec$) and use LVG calculations to estimate density and temperature. We shall also employ the $^{13}$CO($J=1-0$) data where available to place constraints on the physical parameters. This paper is organized as follows: Section 2 describes the observations. Section 3 and 4 show the results and data analysis, respectively. In section 5, we discuss the physical properties of clumps and evolutional sequence of GMCs. Finally, we present a summary in section 6. | \subsection{Dense and compact clumps as candidates for proto-cluster condensations} We have carried out sub-mm $^{12}$CO($J=3-2$) observations of GMCs in the LMC which are most extensive and highly sensitive compared to the previous studies. Six GMCs were selected based on the NANTEN CO survey of the LMC, including 3 Type III GMCs actively forming O stars in addition to 3 Type I/II GMCs which are quiet in O-star formation or cluster formation, although the formation of low to intermediate mass star is not excluded. The spatial resolution of $\sim$ 5 pc and the high sensitivity allowed us to identify 32 molecular clumps in these GMCs and to reveal significant details of the warm and dense molecular gas with $n$(H$_2$) $\sim$ 10$^{3-5}$ cm$^{-3}$ and $T_{\mathrm{kin}}$ $\sim$ 10 -- 200 K. The typical mass of the molecular clumps is large, in the range of 5$\times$10$^3$ -- 2$\times$10$^5$ M$_{\odot}$ with radii of 1 -- 12 pc. Of all of our objects, N 159 No.1 or -W shows the strongest concentration of mass of $\sim$ 7$\times$10$^4$ M$_{\odot}$ within a radius of $\sim$ 5 pc. The masses seem to be larger than those of typical Milky Way GMCs such as those in the eta Car region \citep[e.g.,][]{Yonekura2005}% , although the propeties of these galactic GMCs are based on optically thin C$^{18}$O data. We suggest that these are good candidates for the precursors of rich super clusters like R136 in 30 Dor which includes more than 10$^4$ stars in a small volume with a radius of $\sim$ 1 pc. It is of particular interest to look for even denser gas towards them in higher excitation transitions of the sub-mm region. \subsection{Density and temperature of the clumps and implications} The results of our LVG analysis indicate that clumps are distributed from cool to warm in temperature and from less dense to dense in density. These differences of clump properties in density and temperature show good correspondence with the GMC Types based on the star formation activity, as well as with the \Halpha emission of ionized gas associated with each clump. Clumps in Type III GMCs are warm ($T_{\mathrm{kin}}$ $\sim$ 30 -- 200 K) and are either dense ($n$(H$_2$) $\sim$ 10$^{3.5-5}$ cm$^{-3}$) or less dense ($n$(H$_2$) $\sim$ 10$^3$cm$^{-3}$). Clumps in Type II GMCs are either warm ($T_{\mathrm{kin}}$ $\sim$ 30 -- 200 K) or cool ($T_{\mathrm{kin}}$ $\sim$ 10 -- 30 K) and less dense ($n$(H$_2$) $\sim$ 10$^3$ cm$^{-3}$). Clumps in Type I GMC are cool ($T_{\mathrm{kin}}$ $\sim$ 10 -- 30 K) and less dense ($n$(H$_2$) $\sim$ 10$^3$cm$^{-3}$). The physical parameters of clumps are generally correlated with the star formation activity of GMCs and can perhaps be interpreted in terms of evolutionary effects. Our interpretation is that defferences of clump density and temperature represent an evolutionary sequence of GMCs in terms of density increase leading to star formation; Type I/II GMCs are at a young phase of star formation where density has not yet reached high enough values to cause active massive star formation, and Type III GMCs represent the later phase where the average density is higher, including both high and low density sub-types. The high density clumps in Type III GMCs show high $R_{3-2/1-0,clump}$ of 1.0 -- 1.5 and are associated with strong \Halpha flux while the low density clumps in Type III GMCs show low $R_{3-2/1-0,clump}$ of 0.5 -- 1.0 with weak \Halpha flux. We suggest two alternative ideas to explain % the density difference of the clumps % in Type III GMCs; one is that density is being enhanced by shock compression driven by HII regions and the other is that gravitational condensation of each clump plays a role in the density increase. The former may be difficult because the shock front may occupy a small volume which is likely missed with the present 5 pc beam. It seems thus favorable that the latter scenario is working mainly to enhance density. The present study, which resolved the smaller clumps in GMCs at 5 -- 10 pc scales, indicates that the clumps may have physical properties affected by local properties such as the \Halpha distribution. It should be interesting to investigate the variations among these internal clumps and their relation to star formation. \subsection{FUV Heating of the molecular gas in the LMC} The present findings that the $R_{3-2/1-0,clump}$ is well correlated with \Halpha flux suggests that the heating of molecular gas by far-ultraviolet (FUV) photons may be effective in the LMC where the dust opacity is lower and the FUV intensity is higher than in the Milky Way. The molecular gas in the Milky Way is mainly heated by cosmic ray protons of $\sim$ 100 MeV as discussed by a number of authors, although the surface layers of molecular clouds with small visual extinctions at $A$v $\sim$ a few mag or less may be dominated by the FUV heating \citep[e.g.,][]{Kaufman1999}. Some authors have made detailed calculations of gas heating and cooling under the effects of FUV radiation fileds \citep{Kaufman1999}. We shall try to present a picture that can be applied to the present results below. First, the gas temperature is determined through the balance between the cooling and heating. According to Table 4 in \citet{Goldsmith1978}, the total cooling rate is 6.8$\times$10$^{-27}T^{2.2}$ ergs cm$^{-3}$ s$^{-1}$ for $X$(CO)/(d$v$/d$r$) = 4$\times$10$^{-5}$ and $n$(H$_2$) = 10$^3$ cm$^{-3}$. In 30 Dor region, since $X$(CO)/(d$v$/d$r$) = 3$\times$10$^{-6}$ / 0.8 = 3.75$\times$10$^{-6}$ and this value is 10 times lower than the value used in \citet{Goldsmith1978}, $n$(H$_2$) = 10$^4$ cm$^{-3}$ can be read 10$^3$ cm$^{-3}$, then the cooling rate is estimated as 1.7$\times$10$^{-22}$ ergs cm$^{-3}$ s$^{-1}$ for $T$ = 100 K. This value is a factor 2 -- 3 smaller than that of Galactic clouds with $n$(H$_2$) = 10$^4$ cm$^{-3}$ and $T$ = 50 K. We shall assume that the heating by cosmic ray electrons is not important in the LMC. This assumption is not directly confirmed, but it is consistent with the low non-thermal fraction of the LMC's radio continuun emission \citep[e.g.,][]{Hughes2006} and studies that suggest a significant fraction of cosmic ray electrons are able to escape from low luminosity galaxies \citep[e.g.,][]{Bell2003, Skillman1988} The FUV flux ($G_0$) is estimated as 3500 for 30 Doradus \citep{Bolatto1999, Poglitsch1995, Werner1978, Israel1979}, and 300 for N 159 \citep{Bolatto1999, Israel1996, Israel1979}. In Orion, it is estimated as 25 \citep{Bolatto1999, Stacey1993}. The FUV flux in the LMC is larger than that in the Milky Way. PDR models are calculated by \citet{Kaufman1999} which incorporate the chemical and physical processes that form and destruct atoms or molecules, as well as ionization effects. Figuer 1 of \citet{Kaufman1999} shows the kinetic temperature for a molecular gas layer with density of $n$ (cm$^{-3}$) under FUV flux of $G_0$ at the surface. PDR surface temperatures are estimated as listed in Table \ref{tbl07}. These indicate that temperature becomes as high as 100 -- 300 K on the PDR surface under the conditions of the clumps in Type III GMCs in the LMC. These temperatures are basically consistent with the temperatures of the warm clumps in the present sample. Generally speaking, at a scale of $\sim$ 10 pc, $T_{\mathrm{kin}}$ $\sim$ 100 K seems to be higher than the kinetic temperatures typical in Milky Way GMCs, where the Milky Way values are usually derived from the $^{12}$CO($J=1-0$) emission only \citep[e.g., eta Car $T_{\mathrm{kin}}$ $\sim$ 50 K by][]{Yonekura2005}. This suggests that the heating of molecular clouds may be stronger in the LMC than in the Milky Way and the molecular temperature may be higher. If this is correct, the lower metallicity, resulting in lower extinction, is the basic cause for the higher temperature in addition to the stronger FUV field in the LMC. We shall note in the end that this higher temperature in the molecular gas possibly leads to an increase of the Jeans mass of molecular clumps, which may favor the formation of rich super clusters in the LMC. This is consistent with the higher mass of the molecular clumps which may represent precursors of the clusters. The present work has undertaken to sample 6 GMCs (7 regions) to have a uniform determination of the density and temperature in the LMC. The number of GMCs is still limited to 6 among $\sim$ 300 detected with NANTEN. We should make more efforts to collect appropriate data sets in the sub-mm wavelengths to improve our understanding of the cloud properties. NANTEN2, ASTE and others will certainly be powerful in achieving this goal. | 7 | 10 | 0710.4202 |
0710 | 0710.3600_arXiv.txt | {Two separate statistical tests are described and developed in order to test un-binned data sets for adherence to the power-law form. The first test employs the TP-statistic, a function defined to deviate from zero when the sample deviates from the power-law form, regardless of the value of the power index. The second test employs a likelihood ratio test to reject a power-law background in favor of a model signal distribution with a cut-off. } \begin{document} | The question of whether the cosmic ray energy spectrum exhibits a cut-off at the very highest energies is of central interest to the cosmic ray (CR) physics\cite{refG,refZK}. The flux of CR's at these energies is very small - about $3 /$km$^{2}$ steradian century - and, therefore, statistical analysis techniques which clearly quantify ones knowledge of flux suppression are useful. In this note we apply the statistics first developed for binned CR data sets in \cite{refHague} to an un-binned analysis. We also introduce a new test based on a likelihood ratio test and show that both statistics can quantify our knowledge of a flux suppression. We first establish the mathematical foundations of the analysis. The CR flux follows a power-law for over 10 orders of magnitude. The fundamental probability distribution function (p.d.f.) governing the power-law assumption (normalized such that $\langle \rangle_{X} \equiv \int_{x_{min}}^{\infty} f_{{\scriptscriptstyle X}}(x; x_{min}, \tg)dx=1$) is \begin{equation}\label{equpwlpdf} f_{{\scriptscriptstyle X}}(x; x_{min}, \tg) = A\,x^{-\tg}, \end{equation} where $A=(\tg-1)x^{\tg-1}_{min}$ and the parameter $\tg$ is referred to as the {\it spectral index}. The $n^{th}$ raw moment of this distribution diverges\cite{refNewm} for $n \geq 2$ with $\tg \leq 3$. Alternatively, the expected value of $\ln(x/x_{min})$ is better behaved and offers a crucial result of this analysis. Analytically we find, \begin{equation}\label{equnu} \nu_{n} \equiv \left\langle \ln^n \left( \frac{x}{x_{min}} \right) \right\rangle_{X} = \frac{n!}{(\tg-1)^{n}}. \, \end{equation} For a given sample we use, \begin{equation} \label{equnuhat} {\hat \nu}_{n}(X_{(j)}) \equiv \frac{1}{N-(j-1)} \sum_{i=j}^N \ln^{n} \left( \frac{X_{(i)}}{X_{(j)}} \right). \, \end{equation} In eq.\ref{equnuhat} we denote the sorted (from least to greatest) data set as $\left\{X_{(1)},X_{(2)},\ldots,X_{(N)}\right\}$. To apply these statistics to an un-binned data set we calculate ${\hat \nu}_{n}(X_{(j)})$ for each minimum $X_{(j)}$. We also study a toy p.d.f. which is designed to mimic a power-law up to a certain energy but then exhibit a sharp ``Fermi-Dirac like'' cut-off above that energy\cite{refHague}. We follow the parameterization used in \cite{refPAOicrc318}, \begin{equation}\label{equFDpdf} f_{{\scriptscriptstyle FD}}(x; x_{c}, w_{c}, \tg) = \frac{B\,x^{-\tg}}{1 + \exp \left( \frac{\log x - \log x_{c}}{ w_{c}} \right) }, \end{equation} where $B$ is chosen such that $f_{{\scriptscriptstyle FD}}$ is normalized over the interval $[x_{min},\infty)$, i.e. $\langle \rangle_{{\scriptscriptstyle FD}} = 1$. | We began this note by verifying that the log-binned spectral index estimator has more bias and a larger error than the un-binned (maximum likelihood) estimator. We then detailed two un-binned statistical tests sensitive to flux suppression. We show that both tests show high sensitivity for rejecting the power-law hypothesis in favor of a toy flux suppression model and depend only weakly on the true spectral index. Applying these tests to $3500$ events drawn from a toy cut-off distribution (see eq. \ref{equFDpdf}) we can reject the power-law model in favor of the cut-off model at a confidence level $\sim 4$ standard deviations. | 7 | 10 | 0710.3600 |
0710 | 0710.4485_arXiv.txt | The dissolution process of star clusters is rather intricate for theory. We investigate it in the context of chaotic dynamics. We use the simple Plummer model for the gravitational field of a star cluster and treat the tidal field of the Galaxy within the tidal approximation. That is, a linear approximation of tidal forces from the Galaxy based on epicyclic theory in a rotating reference frame. The Poincar\'e surfaces of section reveal the effect of a Coriolis asymmetry. The system is non-hyperbolic which has important consequences for the dynamics. We calculated the basins of escape with respect to the Lagrangian points $L_1$ and $L_2$. The longest escape times have been measured for initial conditions in the vicinity of the fractal basin boundaries. Furthermore, we computed the chaotic saddle for the system and its stable and unstable manifolds. The chaotic saddle is a fractal structure in phase space which has the form of a Cantor set and introduces chaos into the system. | The dissolution process of star clusters is an old problem in stellar dynamics. Once a star cluster has formed somewhere in a galaxy, it tends to lose mass due to dynamical interactions until it has completely dissolved. It turned out that the physics behind the dissolution of star clusters is intricate and fascinating for theory. If a star cluster of finite mass were isolated, in virial equilibrium (i.e. $\langle v_e^2 \rangle = 12\sigma_{1D}^2$) and the velocity distribution given by a Maxwellian $f_M(X) = \left(4/\sqrt{\pi}\right) X^2 \exp(-X^2)$, where $X=v/(\sqrt{2}\sigma_{\rm 1D})$ and $v, v_e$ and $\sigma_{1D}$ are the velocity, the escape velocity and the velocity dispersion, respectively, the fraction of stars which are faster than the rms escape speed were given by $\chi_e = \int_{\sqrt{6}}^\infty f_M(X) dX = 2\sqrt{6/\pi} \exp(-6) + {\rm erfc}\sqrt{6} \simeq 0.00738316$. This simple analytical result was published by Ambartsumian (1938) and two years later, independantly by Spitzer (1940) who named this effect ``evaporation''. The relevant process which brings stars above the escape speed and lets them evaporate, is two-body relaxation. The time scale of relaxation, which determines the rate of dynamical evolution of a star cluster, yields thus an upper limit to the lifetime of any star cluster. However, since $\chi_e$ is so small (and relaxation time relatively long), the evaporation time is much longer than a Hubble time for typical globular clusters. Following a suggestion of Chandrasekhar (1942), King (1959) studied the effect of ``potential escapers''. These are stars which have been scattered above the escape energy but which have not yet left the cluster. These may be scattered back to negative energies within a crossing time and remain bound. H\'enon (1960, 1969) stressed the importance of few close encounters between stars for the rate of mass loss of star clusters. The Fokker-Planck approximation, which is widely used to study the dynamical evolution of star clusters, neglects strong encounters by construction. Nevertheless, close encounters could still be interpreted statistically as a certain discontinuous Markov process (Tscharnuter 1971). However, direct $N$-body models seem to be ideally suited to study this phenomenon in more detail. Spitzer \& Shapiro (1972) estimate that `` close encounters may produce effects perhaps as great as 10 percent of the ``dominant'' distant encounters'' and ignore them. Nature provides an environment for star clusters in which the escape rate is typically strongly enhanced as compared with the slow evaporation rate of isolated star clusters: The tidal field of a galaxy induces saddle-like troughs in the walls of the potential well of a star cluster (cf. Figure 1). It therefore lowers the energy threshold in star clusters above which stars can escape from zero to a negative value (Wielen 1972, 1974). Moreover, if we consider a star cluster in the tidal field of a galaxy as a dynamical system, the tidal field can change the system's dynamics in a dramatical way as compared with an isolated system. In general, the escape process from star clusters in a tidal field proceeds in two stages: (1) Scattering of stars into the ``escaping phase space'' by two-body encounters and (2) leakage through openings in the equipotential sufaces around saddle points of the potential. The ``escaping phase space'' is defined as the subset of phase space, from which escape is possible. It is well understood that the time scale for a star to complete stage (1) scales with relaxation time. On the other hand, the time scale for a star to complete stage (2) depends mainly on its energy (but also on its location in phase space as we will see). When we neglect the effect of two-body relaxation for the consideration of stage (2), the motion of a single star in the star cluster is determined between times $t_1$ and $t_2$ only by the smooth gravitational potential in which the star moves. The potential itself is generated by the other stars in the star cluster disregarding their ``grainyness'' and by the superposed galactic gravitational field, which is due to the matter distribution of the galaxy. Within this framework, we will study stage (2) of the escape process in this paper. In this connection, the work of Fukushige \& Heggie (2000) is of major interest. Their main result is an expression for the time scale of escape for a star in stage (2) which has just completed stage (1). The dependance of the escape process on two time scales which scale differently with the particle number $N$ imposes a severe scaling problem for $N$-body simulations. The scaling problem is of relevance since the it is on today's general-purpose hardware architectures not yet simply feasible to simulate the evolution of globular clusters with realistic particle numbers of a few hundred thousands or even millions of stars by means of direct $N$-body simulations. The result of Fukushige \& Heggie has been applied in Baumgardt (2001) to solve the important scaling problem for the dissolution time of star clusters in the special case of circular cluster orbits. The obtained scaling law $t_{dis}\propto t_{rh}^{3/4}$, where $t_{dis}$ and $t_{rh}$ are the dissolution and half-mass relaxation times, respectively has been verified, e.g. in Spurzem et al. (2005). The problem of escape has also a long history in the context of the theories of dynamical systems and chaos. It is well-known for a long time, that certain Hamiltonian systems allow for escape of particles towards infinity. Such ``open'' Hamiltonian systems have been studied by Rod (1973), Churchill et al. (1975), Contopoulos (1990), Contopoulos \& Kaufmann (1992), Siopis et al. (1997), Navarro \& Henrard (2001) and Schneider, T\'el \& Neufeld (2002). The related chaotic scattering process, in which a particle approaches a dynamical system from infinity, interacts with the system and leaves it, escaping to infinity, was investigated by many authors, as Eckhardt \& Jung (1986), Jung (1987), Jung \& Scholz (1987), Eckhardt (1987), Jung \& Pott (1989), Bleher, Ott \& Grebogi (1989) and Jung \& Ziemniak (1992). Chaotic scattering in the restricted three-body problem has been studied by Benet et al. (1997, 1999). Typically, the infinity acts as an attractor for an escaping particle, which may escape through different exits in the equipotential surfaces. Thus it is possible to obtain basins of escape (or ``exit'' basins), similar to basins of attraction in dissipative systems or the well-known Newton-Raphson fractals. Special types of basins of attraction (i.e. ``riddled'' or ``intermingled'' basins) have been explored by Ott et al. (1993) and Sommerer \& Ott (1993, 1996). Basins of escape have been studied by Bleher et al. (1988), and they are discussed in Contopoulos (2002). Reasearch on escape from the paradigmatic H\'enon-Heiles system has been done by de Moura \& Letelier (1999), Aguirre, Vallejo \& Sanju\'an (2001), Aguirre \& Sanju\'an (2003), Aguirre, Vallejo \& Sanju\'an (2003), Aguirre (2004) and Seoane Seoane Sep\'ulveda (2007). These papers served as the basis of this work. Relatively early, it was recognized, that the key to the understanding of the the chaotic scattering process is a fractal structure in phase space which has the form of a Cantor set (Cantor 1884) and is called the chaotic saddle. Its skeleton consists of unstable periodic orbits (of any period) which are dense on the chaotic saddle (e.g. Lai 1997) and introduce chaos into the system (e.g. Contopoulos 2002). The properties of chaotic saddles have been investigated by different authors, as Hunt (1996), Lai et al. (1993), Lai (1997) or Motter \& Lai (2001). Both hyperbolic and non-hyperbolic chaotic saddles occur in dynamical systems. In the first case, there are no Kolmogorov-Arnold-Moser (KAM) tori, which means that all periodic orbits are unstable. In the second case, there are both KAM tori and chaotic sets in the phase space (J. C. Vallejo, priv. comm. and e.g. Lai et al. 1993). We note that all of the above references on the chaotic dynamics are exemplary rather than exhaustive since there exists a vast amount of literature on these topics. The aim of this paper is to allude to the importance of this last-mentioned branch of research for the field of stellar dynamics of star clusters. We will study the escape process from star clusters within the framework of chaotic dynamics. In section 2, we introduce the tidal approximation, i.e. approximate equations of motion for stellar orbits in a star cluster which is embedded in the tidal field of a galaxy. In section 3, we describe our (very simple) model of the gravitational potential. In section 4, we discuss Poincar\'e surfaces of section, which show the effect of a Coriolis asymmetry. Furthermore, we discuss the basins of escape in Section 5 and the chaotic saddle and its stable and unstable invariant manifolds in Section 6. Section 7 contains the discussion and conclusions. | We studied the chaotic dynamics within a star cluster which is embedded in the tidal field of a galaxy. We calculated within the framework of the tidal approximation Poincar\'e surfaces of section, the basins of escape, the chaotic saddle and its stable and unstable manifolds as well. The system is non-hyperbolic which has important consequences for the dynamics, i.e. there are orbits which do not escape if relaxation is neglected. These are mainly the retrograde orbits as has been shown earlier by Fukushige \& Heggie (2000) and, more recently, in the $N$-body study by Ernst et al. (2007). The escape times are longest for initial conditions near the fractal basin boundaries. The decay law is a power law for those stars which escape from the regions without sensitive dependance on the initial conditions in Figure \ref{fig:basins} (i.e. with short escape times, as can be seen in Figure \ref{fig:times}). On the other hand, the decay law is exponential for orbits which escape from the regions with sensitive dependance on the initial conditions (i.e. with long escape times). The effect of relaxation (i.e. a diffusion in the Jacobian and the third integral among different stellar orbits) on the chaotic dynamics which we investigated in this work may be a very interesting topic for future research. | 7 | 10 | 0710.4485 |
0710 | 0710.1096_arXiv.txt | We report first results of a multifrequency campaign from radio to hard X-ray energies of the prominent \gray\ blazar 3C~279, which was organised around an INTEGRAL ToO observation in January 2006, and triggered on its optical state. The variable blazar was observed at an intermediate optical state, and a well-covered multifrequency spectrum from radio to hard X-ray energies could be derived. The SED shows the typical two-hump shape, the signature of non-thermal synchrotron and inverse-Compton (IC) emission from a relativistic jet. By the significant exposure times of INTEGRAL and Chandra, the IC spectrum (0.3 - 100 keV) was most accurately measured, showing -- for the first time -- a possible bending. A comparison of this 2006 SED to the one observed in 2003, also centered on an INTEGRAL observation, during an optical low-state, reveals the surprising fact that -- despite a significant change at the high-energy synchrotron emission (near-IR/optical/UV) -- the rest of the SED remains unchanged. In particular, the low-energy IC emission (X- and hard X-ray energies) remains the same as in 2003, proving that the two emission components do not vary simultaneously, and provides strong constraints on the modelling of the overall emission of 3C~279. | The discovery by the experiments aboard the Compton Gamma-Ray Observatory (CGRO) that blazars can radiate a large -- sometimes even the major -- fraction of their luminosity at \gray\ energies marked a milestone in our knowledge on blazars. During the CGRO mission about 90 blazars were detected by the different CGRO experiments at \gray\ energies, the majority by the EGRET experiment at energies above $\sim$100~MeV \citep{Hartman99}. 3C~279, an optically violently variable (OVV) quasar located at a redshift of 0.538, is one of the most prominent representatives of these sources. The source shows rapid variability in all wavelength bands, polarized emission in radio and optical, superluminal motion, and a compact radio core with a flat radio spectrum. These properties put the quasar 3C~279 into the blazar sub-class of AGN. According to the unified model of Active Galactic Nuclei (AGN), blazars are sources which expel jets close to our line-of-sight. 3C~279 was already detected by INTEGRAL in June 2003 \citep{Collmar04}. Those high-energy observations were supplemented in X-rays by a short (5 ksec) Chandra pointing, and by ground based monitoring from radio to optical bands. Since the blazar was found in 2003 at the faintest optical brightness (optical R-band: $\sim$17~mag) of the last 10 years, roughly 5 mag fainter than the maximum, and about 2.5 to 3 mag fainter than average, a simultaneous spectral energy distribution (SED) of an exceptional optical low-state could be compiled \citep{Collmar04}. In order to measure emission changes as function of optical brigthness, we proposed for an INTEGRAL ToO observation during an optical high state (criterion: optical R-band brigther than 14.5~mag). In January 2006, the trigger citerion was met, and the INTEGRAL observations together with supplementing multifrequency observations were carried out. This campaign resulted in a well covered SED from radio to hard X-ray energies. In this paper we present first results of this 2006 multiwavelength campaign on 3C~279. Because of the page restrictions, we concentrate on presenting the main observational results, focussing on the observed SED and its comparison to the one of 2003. A more detailed presentation, including a discussion on the scientific implications of the new results, will be given in a later paper (Collmar et al., in prep.; including also all participating individual WEBT (Whole Earth Blazar Telescope) collaborators in the author list). In addition, the results on variability analyses and time correlations of the different wavelength bands will be given elsewhere (B\"ottcher et al., in prep.). \begin{figure}[th] \centering \epsfig{figure=fig_1a.eps,width=8.0cm,clip=} \epsfig{figure=fig_1b.eps,width=8.1cm,clip=} \caption{ Top: The ISGRI image shows a 6.4\sig\ detection of 3C~279 in the 20-60~keV band. In addition, the Seyfert galaxy NGC 4593 is even more clearly detected. \newline Bottom: The ISGRI hard X-ray spectrum between 20 and 100 keV is shown, together with the best-fit power-law shape. \label{fig:1} } \end{figure} | 7 | 10 | 0710.1096 |
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0710 | 0710.2789_arXiv.txt | {In December 2004, the soft gamma-ray repeater \src\ emitted the most powerful giant flare ever observed. This probably involved a large-scale rearrangement of the magnetosphere leading to observable variations in the properties of its X-ray emission. Here we present the results of the first \suz\ observation of \src, together with almost simultaneous observations with \xmm\ and \int. The source seems to have reached a state characterized by a flux close to the pre-flare level and by a relatively soft spectrum. Despite this, \src\ remained quite active also after the giant flare, allowing us to study several short bursts observed by \suz\ in the 1--100 keV range. We discuss the broad-band spectral properties of \src, covering both persistent and bursting emission, in the context of the magnetar model, and consider its recent theoretical developments. | The four known Soft Gamma-ray Repeaters (SGRs) were discovered as transient sources of high-energy photons; they emit sporadic and short ($\sim$0.1 s) bursts of (relatively) soft gamma-rays with luminosity $L\sim10^{40}$--$10^{41}$ \lum during periods of activity, that are often broken by long intervals of quiescence. Three ``giant'' flares with luminosity \mbox{$\gtrsim$$10^{43}$ \lum} have also been observed to date, each one from a different SGR: on March 5, 1979 from SGR\,0526--66 in the Large Magellanic Cloud \citep{mazets79}, on August 27, 1998 from SGR\,1900+14 \citep{hurley99}, and on December 27, 2004 from \src\ \citep{hurley05}. Persistent emission with $L\sim10^{35}$ \lum\ is also observed from SGRs in the soft X--ray range (\mbox{$<$10 keV}) and, for \src\ and SGR\,1900+14, also in the hard X-ray range \citep{mgm05,gotz06}. In three cases, periodic pulsations at a few seconds have been detected. The bursts, the giant flares, the quiescent X-ray counterparts, and the pulsations have been interpreted in the framework of the magnetar model \citep[see][and references therein]{tlk02}. Magnetars are highly magnetized neutron stars with field strengths of \mbox{$10^{14}$--$10^{15}$ G}, larger than those of the majority of radio pulsars. The ultimate source of energy for the bursts and the quiescent emission is believed to be the ultra-strong magnetic field.\\ \indent \src\ was discovered in 1979 \citep{laros86,laros87} and its persistent X-ray counterpart was observed for the first time with the \asca\ satellite in 1993 \citep{murakami94}. A \xte\ observation led to the discovery of pulsations in the persistent emission with period $P\simeq7.47$ s and period derivative $\dot{P}\simeq2.6\times10^{-3}$ s yr$^{-1}$ \citep{kouveliotou98}. Under the assumption of pure magnetic dipole braking, these values imply a surface magnetic field strength of \mbox{$8\times10^{14}$ G}, strongly supporting the magnetar model. Both the burst rate and the X-ray persistent emission of \src\ started increasing during 2003 and throughout 2004 \citep{mte05,tiengo05,met07,woods07}, culminating with the giant flare of December 27, 2004, during which $\sim$$10^{47}$ erg were released\footnote{Assuming isotropic luminosity and for a distance $d=15$ kpc \citep{corbel97,mcclure05}. } \citep{hurley05,mereghetti05,terasawa05}. This giant flare was exceptionally intense, $\sim$100 times more energetic than those from SGR\,0526--66 and SGR\,1900+14. Observations with \xte\ unveiled, for the first time in an isolated neutron star, rapid quasi-periodic oscillations in the pulsating tail of the flare, likely related to global seismic oscillations on the neutron star surface \citep{ibs05}. The flare produced a hard X-ray (\mbox{$>$80 keV}) afterglow lasting a few hours \citep{mereghetti05,frederiks07} and a radio afterglow that faded in a few days \citep{cameron05}. The small positional uncertainty of the radio observations permitted to identify the likely IR counterpart of the SGR \citep{kosugi05,israel05}. The fluxes observed in the IR and gamma energy bands show a variability correlated with that observed in the 2--10 keV energy range \citep{met07}.\\ \indent After the giant flare, the persistent X-ray flux of \src\ started to decrease from its outburst level, and its X-ray spectrum to soften \citep{rea05,rea05_atel,met07,tiengo05,woods07}. A Similar flux decrease have been observed from its radio afterglow \citep{gaensler05,taylor05} and its newly discovered IR counterpart \citep{israel05,rea05_atel,met07}.\\ \indent Here we present the results of the first \suz\ observation of \src, covering both persistent and bursting emission in the 1--100 keV energy band. We also report on the analysis of a simultaneous observation performed with \xmm\ and the latest outcomes of the monitoring of \src\ with \int, comparing them with what is seen in the same energy ranges with \suz. | The \suz, \xmm, and \int\ observations reported here represent a complementary data set that allows us to study the spectral properties of \src\ in the broad \mbox{1--100 keV} energy range. Although the \suz/HXD does not have imaging capabilities, we know thanks to \int\ that no other bright hard X-ray sources are present in its field of view. The uncertainties in the instrumental background (currently at the $\sim$5\% level) and in the modelling of the Galactic Ridge emission are a more relevant concern. Future improvements in the knowledge of these components may eventually allow us to obtain more robust conclusions. Thanks to its imaging capabilities, background issues do not affect the \int\ observations. However, the IBIS/ISGRI data required long integration times, with discontinuous observations spanning several months. Thus they provide information on the average properties only. The possible presence of long term variability was in fact one of the main motivations to perform the \suz\ and \xmm\ observations simultaneously. With all these caveats in mind we proceed now to discuss the broad band spectrum of \src.\\ \indent The \xmm\ and \int\ spectra are consistent with an extrapolation of the power-law plus blackbody model measured in the 2--10 keV band. Between 12 and 30 keV the \suz/HXD sensitivity is better than that of \int, allowing to detect \src\ during a single 50 ks long observation. With respect to the \xmm\ and \int\ joint fit, the HXD data show an ``excess'' (see Fig.\,\ref{broad}) that cannot be completely ascribed to calibration uncertainties between the various instruments.\\ \indent Given the lack of a direct measure of the Galactic Ridge emission around \src, we cannot exclude that this excess is due to an underestimation of such contribution to the background. If instead the excess is a real feature of the spectrum of \src, its broadband spectrum could be empirically modeled adopting a power-law with the photon-index changing from $\sim$1 to $\sim$2 at $\sim$16 keV, and a blackbody component with $k_BT\sim0.8$ keV. This would agree with the results reported by \citet{gotz06}, who point out that the hard tails of the SGRs are softer than the power-law components measured below \mbox{10 keV}.\\ \indent The presence of a down-break in the 10--20 keV spectrum of \src\ would have remarkable physical implications. The soft X-ray emission from magnetar candidates (SGRs and AXPs) is usually interpreted within the twisted magnetosphere model as due to resonant cyclotron up-scattering of soft photons from the star surface onto charges flowing into the magnetosphere \citep{tlk02}. Detailed calculations of resonant compton scattering (RCS) spectra have been recently performed \citep{lyutikov06,fernandez07} and successfully applied to fit AXP spectra \citep{rzl07}. Quite interestingly, some of the model spectra presented by \citet{fernandez07} exhibit a downward break in the tens of keV range. Their overall shape is quite reminiscent of the \suz\ XIS/HXD-PIN spectrum of \src, and, as noted by \citet{fernandez07}, they may also play a role in the interpretation of the broadband X-ray spectrum of SGR\,1900+14. In particular, when assuming a (non-thermal) top-hat or a broadband velocity distribution for the magnetospheric charges, multiple peaks can appear in the spectrum (see their Fig.\,6 and Fig.\,11). The downturn possibly present in our data may then be due the presence of a second ``hump'' (in addition to the main thermal one) in the range \mbox{10--20~keV}. Nobili, Turolla, \& Zane (in preparation) assuming a 1-D thermal electron distribution superimposed to a (constant) bulk velocity, found also double humped spectra. In this case the second (and only) hump occurs when resonant scattering is efficient enough to fill the Wien peak at the temperature of the comptonising particles. A spectral break at $\sim$15~keV would translate then in a temperature of $\sim$5 keV for the magnetospheric electrons. If a more refined treatment of background subtraction confirms the spectral break in the X-ray data of \src\, this would provide important diagnostics for the physical parameters of the model.\\ \indent In 2003 we started a long-term monitoring program to study the time evolution of the spectral properties of \src\ using the \xmm\ X-ray satellite. The December 27, 2004 giant flare was a fortunate occurrence that allowed us to observe how the source properties evolved in the two years leading up to the flare and how they changed after this dramatic event \citep[see][for details]{mte05,met07,rea05_atel,tiengo05}. The \xmm\ data showed a doubling of the flux in the September--October 2004 followed by a gradual recovery to the ``historical'' level after the giant flare. A direct comparison of the \xmm/pn count rates measured in the different observations shows that before the giant flare the flux of \src\ in the \mbox{1--10 keV} band was monotonically increasing, while the three observations after the flare, and preceding the one reported here, followed a steady decreasing trend. The September 2006 observation breaks this long term decay, having a count rate higher by 5\% with respect to the last \xmm\ observation performed 5 months before. This slight (but statistically significant) re-brightening might indicate either a temporary oscillation around an equilibrium flux level or the start of a new monotonic flux increase, similar to the one that preceded the December 2004 giant flare. This phenomenon was interpreted as due to the building up of a magnetospheric twisting, that determined also the hardening of the X-ray spectrum, an increase of the spin-down and a more intense bursting activity \citep{mte05}. The relatively high burst rate observed during the \xmm\ and \suz\ observations of September 2006 (see Section\,\ref{burstingemission}) is therefore another indication of a possible increase of the magnetospheric twisting in \src, but, before a new \xmm\ observation will be performed, only the monitoring of the frequency and intensity of \src\ bursts can tell us if the evolution is erratic or follows a stable trend. The recent report of a bright burst from \src\ \citep{golenetskii07,perotti07} seems actually to favour the second hypothesis.\\ \indent Two of the bursts detected during the \suz\ observation were bright enough to allow spectral analysis. In both cases, the broadband spectrum (\mbox{2--100 keV}) revealed the presence of two components: a soft component which is well reproduced by a blackbody with $k_BT \sim 2$--4~keV, and a harder one whose spectral shape is not firmly established and can be equally well fit with a power-law, a hot bremsstrahlung or a second blackbody (See Table\,\ref{bfits} and Fig.\,\ref{burstfits}). In absence of robust theoretical predictions, we can not exclude that a two component model simply reflects our ignorance of the correct spectral shape, and has therefore a purely phenomenological significance. However, it is worth noticing that, from their recent analysis of a sample of 50 bursts detected from \src\ with \hete\ from 2001 to 2005, \citet{nakagawa07} have suggested the presence of a time delay between the 30--100~keV and the 2--10~keV emission. Although such a delay can be attributed to an intrinsic, rapid spectral softening, an alternative, and simpler, interpretation invokes the presence of two separate emitting regions.\\ \indent Let us consider a scenario in which the two components are physically distinct and let us consider the hard component first. In the magnetar scenario, short bursts are usually ascribed to either reconnection phenomena in the external magnetosphere \citep[eventually modulated by a tearing instability, see][]{lyutikov03} or movements of the footprints of the external magnetic field, produced by crustal deformations or fractures driven by the stress exerted by the internal field helicity \citep{thompson95,thompson01,tlk02}. Both kind of processes lead to the generation and launch of an Alfv\`en wave, which produces and accelerates a particle cascade, and ultimately is detected as a burst. The emerging spectrum is expected to be synchrotron dominated, unless the Alfv\`en wave is temporarily trapped in a fireball and thermalized. Therefore, both the \mbox{BB + PL} and BB + BB spectral fits are consistent with a scenario in which such an Alfv\`en wave is responsible for the hard component. We notice that, although a fireball formation is not required to explain short bursts (and therefore usually not invoked in such cases), to our knowledge there is no a priori reasons why a small fireball can not be created and evaporated in a sub-second time scale, giving rise to a thermal spectrum. A point in favour of this interpretation is that, in the BB + BB fit, the temperature of the hot blackbody is remarkably close to the minimal temperature above which a fireball thermalizes, $k_BT \approx 11\,(R_{\rm{NS}}/10\,\rm{km})^{-1/5}$~keV, according to \citet{thompson01}, eq. [71] \citep[see also][for similar findings based on longer duration burst]{olive04}. The third model of the hard component which is compatible with our data is a bremsstrahlung emission at $\sim$100~keV. Quite recently, \citet{thompson05} and \citet{beloborodov07} discussed the electrodynamics of the magnetar coronae and the production mechanisms for soft gamma-rays. In particular, their model predicts the existence of a thin transition layer between the corona and the thermal photosphere, where Langmuir turbulence can be excited by a downward beam of current-carrying charges. As a result, the transition layer can be heated up to a typical temperature of $\sim$100~keV, and emit, approximatively, an optically thin bremsstrahlung at a single temperature. Although \cite{thompson05} model was originally developed in connection with the persistent hard emission of magnetars, the predicted bremsstrahlung temperatures are remarkably close to those we detected during the two bursts, suggesting that a similar mechanism may instead be activated during periods of activity.\\ \indent Our results about the spectral modelling of the soft X-ray component are more robust, inasmuch as all our spectral fits require the presence of a cold blackbody with \mbox{$k_BT \sim 2$--4~keV}. This is in agreement with similar findings by \cite{olive04}, \cite{nakagawa07}, and \citet{feroci04}. This component is usually interpreted as due to emission from a fraction of the star surface (which can be as large as the whole star in the case of our BB + BB fit) heated by returning currents. Alternatively, it has been suggested that the soft component may originate up in the magnetosphere ($\leq$700~km), presumably due to a delayed emission process \citep{nakagawa07}. Here we only notice that, although the spectra of our two events are compatible with emission from the star surface, the radius of the cold blackbody as measured during other short bursts can reach values much higher than 50--100~km \citep[][similar findings have been found in the case of SGR\,1900+14 bursts measured with \emph{Swift}, Israel et al. (in preparation)]{nakagawa07}. One possible explanation is that part of the flare energy is intercepted and reprocessed in a larger region and re-emitted at a lower temperature. In such scenario, the radius of the reprocessing region can then vary depending on the fraction of material that is intercepted, and is not bounded by the value of the star radius. \citet{thompson01} considered the equilibrium state of a pair corona sufficiently extended that the local value of the magnetic field is $B\ll B_{\rm{QED}}$, so that photon splitting can be ignored (if the magnetic field scales as a dipole, this occurs above $\sim$3 star radii, for a polar surface value of $\sim$$10^{15}$~G). They found that a stable balance between heating and diffusive cooling requires a continuous source of ordinary photons that can be provided, for instance, by external illumination. If the corona intercepts part of the flare beam (which in their treatment is assumed to originate in a trapped fireball, although in general this is not necessarily required), equilibrium is possible below a critical luminosity given by \begin{displaymath} L<L_{\rm{max}} = 1.5 \times 10^{42}\, \tau_{EO}^{-1} \left ( \frac {k_BT_e} {20\ {\rm keV} } \right )^4 \left ( \frac {R }{10\ {\rm km}} \right )^2 \, {\rm erg\ s^{-1}} \, , \end{displaymath} where $T_e$ is the pair temperature in the corona, $\tau_{EO} \sim 1$ is the scattering depth for ordinary to extraordinary mode conversion and $R$ the radius of the emitting part of the corona (see equations [84] and [89] in \cite{thompson01} and note that there is a typo in their equation [89]: it should contain $R^2$ instead of $R^{-2}$). Even for an emission region as small as 5--20~km (as inferred by our best fit of the low temperature blackbody) and a temperature of $k_BT \sim 2$--4~keV, this is well above the luminosity emitted during the two bursts detected by \suz. Therefore, simply on the basis of energetics, relatively large emitting regions for the cold blackbody are compatible with the temporary formation of a pair corona, sustained by a fraction of the flare energy. When the heating rate ceases, the pair atmosphere contracts and quickly evaporates. In order to derive firmer conclusions a more detailed analysis is needed, mainly in assessing the possibility that the intercepted beam is thermalized and re-emitted as a blackbody. This study is beyond the purpose of this paper, and will be presented elsewhere (Israel et al. 2007, in preparation). | 7 | 10 | 0710.2789 |
0710 | 0710.0090_arXiv.txt | An $\sim$800~arcmin$^2$ mosaic image of the W3 star forming complex obtained with the {\it Chandra X-ray Observatory} gives a valuable new view of the spatial structure of its young stellar populations. The {\it Chandra} image reveals $\sim 1300$ faint X-ray sources, most of which are PMS stars in the cloud. Some, but not all, of the high-mass stars producing hypercompact and ultracompact H{\sc II} (UCHII) regions are also seen, as reported in a previous study. The {\it Chandra} images reveal three dramatically different embedded stellar populations. The W3~Main cluster extends over 7~pc with $\sim 900$ X-ray stars in a nearly-spherical distribution centered on the well-studied UCHII regions and high-mass protostars. The cluster surrounding the prototypical UCHII region W3(OH) shows a much smaller ($\leq 0.6$~pc), asymmetrical, and clumpy distribution of $\sim 50$ PMS stars. The massive star ionizing the W3~North H{\sc II} region is completely isolated without any accompanying PMS stars. In W3~Main, the inferred ages of the widely distributed PMS stars are significantly older than the inferred ages of the central OB stars illuminating the UCHIIs. We suggest that different formation mechanisms are necessary to explain the diversity of the W3 stellar populations: cluster-wide gravitational collapse with delayed OB star formation in W3~Main, collect-and-collapse triggering by shock fronts in W3(OH), and a runaway O star or isolated massive star formation in W3~North. | } W3 (Westerhout 3) is perhaps the most active region of current star formation in the nearby Galaxy. Extending 30~pc along the edge of a $M \simeq 5 \times 10^4$~M$_\odot$ giant molecular cloud (GMC), the star-forming complex has dozens of embedded young massive stars producing a variety of pre-stellar condensations, hot molecular cores, hypercompact to small H{\sc II} regions, maser clusters, and molecular outflows \citep[e.g.,][]{Lada78, Reid80, Dreher81, Tieftrunk97, Chen06}. Its infrared sources have an integrated luminosity of several~$\times~ 10^5$~L$_\odot$. Situated just east of W3 are the older IC~1795 and IC~1805 clusters, the latter lying within the enormous W4 superbubble/chimney structure blown by generations of massive stars. The W4--IC~1795--W3 complex is widely considered to be an examplar of sequential triggered star formation \citep{Lada78, Oey05}. Recent SCUBA observations of the W3 GMC find a higher percentage of the gas mass gathered into dense molecular clumps at the eastern edge compared to the undisturbed parts of the W3 GMC, supporting this triggering scenario \citep{Moore07}. A detailed description of the W3 and W4 complexes and a thorough review of the literature are given by \citet{Megeath07}. The richest site of massive star formation in W3 is the W3~Main cluster of embedded OB stars, dominated by the very young and luminous IRS~4 and IRS~5 sources. IRS~5 lies at the center of a 0.1~pc concentration of massive stars resembling a nascent counterpart of the Orion Trapezium \citep{Megeath05}. W3(OH) to the southeast and W3~North to the north have massive stars but appear less active than W3~Main. The distance to the complex is accurately measured from maser kinematics to be 2.0~kpc \citep{Xu06, Hachisuka06}. Despite the intense study of W3 at radio, millimeter, and infrared wavelengths, little is known about its low mass stellar population. For example, a $JHK$ near-infrared (NIR) survey of $5\arcmin \times 5\arcmin$ in W3~Main reveals $\sim 40$ sources with $K$-band excesses indicative of Class~I--II pre-main sequence (PMS) stars with disks \citep{Ojha04}. Hundreds of other stars are detected, but infrared photometry cannot discriminate disk-free Class~III PMS stars from the strongly contaminating population of unrelated Galactic field stars (mostly red giants). A new mid-infrared (MIR) photometric survey of W3~Main, IC~1795, and W3(OH) with the {\it Spitzer Space Telescope} helps to identify cluster members \citep{Ruch07}, but it suffers confusion from three effects: foreground and background Galactic field stars \citep{Benjamin05}, bright diffuse emission produced by heated dust around the H{\sc II} regions \citep{Povich07}, and extragalactic objects with MIR excesses \citep{Harvey07}. X-ray surveys of young stellar clusters (YSCs) with the {\it Chandra X-ray Observatory} are surprisingly efficient at detecting low mass PMS populations, even at distances around 2~kpc and at obscurations $10<A_V<150$ mag typical for W3 stars. PMS X-ray emission arises primarily from violent magnetic reconnection events, similar to solar flares but far more powerful, and is largely independent of circumstellar disks or accretion \citep[see reviews by][]{Feigelson99, Feigelson07}. Luminous and spectrally hard X-ray flares are present throughout the PMS phases of Class~I--II--III at levels $10^2-10^3$ above that seen in old disk populations \citep{Preibisch05a}, so relatively few Galactic disk interlopers appear in X-ray samples. These field star X-ray sources and extragalactic contaminants are easily removed \citep{Getman06}. Due to a poorly-understood statistical association between X-ray luminosity and PMS stellar mass \citep{Preibisch05b,Guedel07}, a flux-limited X-ray observation of a young stellar cluster will be roughly complete down to a corresponding mass limit. Taking together these properties of X-ray studies, we find that X-ray surveys at sufficiently high spatial resolution and sensitivity provide uniquely rich, largely disk-unbiased, mass-limited, and nearly uncontaminated samples of PMS stars in both embedded and unobscured YSCs. These samples complement MIR surveys obtained with the {\it Spitzer Space Telescope}, which generally extend down to lower masses (including the brown dwarf regime) but cannot readily discriminate disk-free PMS stars from field stars. {\it Spitzer} thus detects more disky Class~0-I-II systems while {\it Chandra} effectively samples Class~III systems in addition to many Class I and II stars. The X-ray samples are useful for various astrophysical purposes such as probing the stellar Initial Mass Function, protoplanetary disk evolution, and magnetic activity. In an early {\it Chandra} study, two $\sim$20~ks exposures of a $\sim 300$~arcmin$^2$ field in W3~Main revealed 236 X-ray sources \citep{Hofner02}. Several are associated with massive stars ionizing H{\sc II} regions but most do not have counterparts in $JHK$ images. We report here an extension of those efforts with a {\it Chandra} mosaic of 7 exposures totaling $\sim 230$~ks over $\sim 800$~arcmin$^2$, spanning much of the W3 star forming complex (Figure~\ref{fig:ACIS_mosaic}). A preliminary discussion of this mosaic, a {\it Chandra} Large Project, is given by \citet{Townsley06}. Over 1300 X-ray sources are seen; a full listing and study of their properties will be presented in a separate paper. For each source, {\it Chandra} observations provide a sub-arcsecond position, line-of-sight absorption, and rough mass estimate in addition to magnetic activity characteristics. We discuss here insights into the global structure and origins of the W3 stellar populations derived from the new {\it Chandra} data. The brief presentation of the observations in \S\ref{sec:obs} will be expanded in a forthcoming paper with complete source lists similar to our group's recent studies of the Cep~OB3b \citep{Getman06}, Pismis~24 \citep{Wang07a}, M~17 \citep{Broos07}, RCW~49 \citep{Tsujimoto07}, and Rosette star forming region \citep{Wang07b, Wang07c} YSCs. The three well-studied star forming regions in W3 are described and contrasted in \S\ref{sec:Xray}, and explanations for their origin are considered in \S\ref{sec:diversity}. Section~\ref{sec:W3Main.origin} considers in more detail the implications of the W3~Main results for astrophysical models of star cluster formation. | This paper introduces a new high-resolution X-ray mosaic of the W3 star forming complex, a Large Project of the {\it Chandra X-ray Observatory}. A rich population of $\sim 1300$ young stars is imaged and the three well-known regions of high-mass star formation are shown to have very different populations of low mass stars: W3~Main is a large, rich, nearly spherical cluster; W3(OH) lies in an elongated group of sparse stellar clumps; and W3~North is an isolated O star without low-mass companions. Suggestions of these differences were inferred from earlier infrared studies, but they are more apparent here because the X-ray selection has the advantage of low contamination by the Galactic field population or diffuse interstellar emission, high penetration into molecular environments, and little bias towards stars with massive protoplanetary disks. We emerge from this study with an improved view of star formation in the region. The W3(OH) structures are consistent with collect-and-collapse triggering process caused by by shocks from the older IC~1795 cluster, as previously suggested. The W3~Main cluster, however, does not show the elongated and patchy structure of a recently triggered star cluster and appears to have formed in an earlier episode. Its PMS population strongly resembles those seen in other {\em Chandra} studies of massive star-forming regions such as those ionizing the Orion, M~17 and Rosette Nebulae. A major difference is that the individual H{\sc II} regions in these other clusters have already merged into a large blister and dispersed their natal clouds. In contrast, the W3~Main OB stars are very recently formed with small individual H{\sc II} regions still embedded in a dense, clumpy molecular medium. Star formation in W3 has proceeded in a prolonged fashion, and apparently with a time-dependent Initial Mass Function. The OB stars exciting the hypercompact and ultracompact HII regions at the center of W3~Main formed more recently than the hundreds of X-ray emitting PMS stars distributed over several parsecs. W3~Main thus becomes a critical testbed for theories of rich cluster formation. | 7 | 10 | 0710.0090 |
0710 | 0710.5026_arXiv.txt | Detection and study of gravitational waves from astrophysical sources is a major goal of current astrophysics. Ground-based laser-interferometer systems such as LIGO and VIRGO are sensitive to gravitational waves with frequencies of order 100 Hz, whereas space-based systems such as LISA are sensitive in the millihertz regime. Precise timing observations of a sample of millisecond pulsars widely distributed on the sky have the potential to detect gravitational waves at nanohertz frequencies. Potential sources of such waves include binary super-massive black holes in the cores of galaxies, relic radiation from the inflationary era and oscillations of cosmic strings. The Parkes Pulsar Timing Array (PPTA) is an implementation of such a system in which 20 millisecond pulsars have been observed using the Parkes radio telescope at three frequencies at intervals of two -- three weeks for more than two years. Analysis of these data has been used to limit the gravitational wave background in our Galaxy and to constrain some models for its generation. The data have also been used to investigate fluctuations in the interstellar and Solar-wind electron density and have the potential to investigate the stability of terrestrial time standards and the accuracy of solar-system ephemerides. | The existence of gravitational radiation is a key prediction of relativistic theories of gravity. These waves propagate at the speed of light and are generated by the acceleration of massive bodies. Astrophysical sources of gravitational waves (GW) include relic radiation from the inflation era \citep{tur97,gri05}, radiation from reconnection and oscillations of cosmic strings \citep{dv05}, supernovae and formation of compact stars and black holes \citep{bsr+05}, binary super-massive black holes in the cores of galaxies \citep{jb03,wl03a,eins04}, coalescence of double-neutron-star binary systems \citep{kkl+04a} and short-period X-ray binaries in our Galaxy \citep{nyp04}. Some of these sources may be individually detectable, others combine to form an essentially isotropic and stochastic background of GW which permeates all of space. Observations of orbital decay of double-neutron-star binary systems have provided irrefutable evidence that gravitational radiation exists and that its power is accurately described by Einstein's general theory of relativity \citep{wt05,ksm+06}. However the signal expected at the Earth from any realistic source is exceedingly weak, with typical strain amplitudes of order $10^{-22}$ at frequencies of order 1 Hz. Despite considerable efforts over more than 40 years, up to now there has been no confirmed direct detection of GW. Initial efforts used the massive bar detectors pioneered by Joseph Weber \citep{web69} but more recent detectors with higher sensitivity are based on laser interferometer systems, for example, the ground-based systems LIGO \citep{aad+92} and VIRGO \citep{gb02} and the proposed space interferometer LISA \citep{dan00}. The ground-based interferometers are sensitive to GW with frequencies in the range 10 -- 500 Hz, whereas LISA is sensitive to frequencies in the range 0.1 -- 100 mHz. Initial LIGO is now operating and has set limits on various sources \citep[e.g.,][]{aaa+04}; higher sensitivity will be achieved with Advanced LIGO which is due for completion in 2011. The launch date for LISA is rather uncertain but is unlikely to be before 2017. Pulsars, especially millisecond pulsars (MSPs) are incredibly precise clocks making possible many interesting applications. Of most interest to us here is the use of MSPs as GW detectors. GW passing over pulsars and over the Earth will modulate the received pulsar period; the net effect is the difference in the modulation at the two ends of the path \citep{det79}. Pulsar timing experiments measure variations of pulse phase relative to model predictions. They are therefore most sensitive to long-period GW with periods comparable to the data span, typically several years, which corresponds to frequencies in the nanoHertz regime. Even for these long periods, the expected timing residuals are very small. Simulations using {\sc tempo2} \citep{hem06,hjl+07} show that a binary system consisting of two $10^9\;\Msun$ black holes with a 4-yr orbital period in a galaxy at redshift 0.5 will produce a timing residual of amplitude just 1.5 ns. Although such a signal would be very difficult to detect with current technology, expected levels of the stochastic GW background from binary super-massive black holes in galaxies are considerably higher with millions of galaxies throughout the Universe contributing. Predicted levels of the GW background from other sources such as the inflation era and cosmic strings, while much more uncertain, are at comparable levels and are potentially detectable. Other sources of ``noise'' exist in pulsar timing data and if we wish to detect GW using pulsar timing we have to be able to separate these different effects. With timing observations of just one or even a few pulsars, upper limits may be set but a positive detection is not possible. However, a large sample of pulsars widely distributed on the sky --- a pulsar timing array --- can in principle {\em detect} GW with frequencies in the nanoHertz range. | Pulsar timing arrays exploit the remarkable stability of MSP periods to enable investigation of a range of phenomena. Direct detection of gravitational waves from astrophysical sources is a major goal of current astrophysics and pulsar timing arrays have the potential to achieve this goal. They are sensitive to GW at frequencies of a few nanoHertz, complementing ground-based and space-based laser interferometer systems which are sensitive at much higher frequencies. PTA systems also have the potential to establish a ``pulsar timescale'' which is more stable than the best terrestrial timescales over intervals of several years or more and to detect errors or omissions in models of Solar-system dynamics, for example, the existence of currently unknown trans-Neptunian objects. The Parkes Pulsar Timing Array (PPTA) project is using the Parkes 64-m radio telescope to time a sample of 20 millisecond pulsars at three frequencies every 2 -- 3 weeks. Observations commenced in early 2005, so we now have over two years of timing data. Sub-microsecond timing residuals have been achieved on about half the sample but we still need to improve timing precisions by a factor of a few in order to have a realistic chance of detecting the stochastic GW background. New instrumentation and other improvements will help us to achieve that goal. We are also actively seeking international collaborations with other timing array projects to increase the sky coverage and the density of observations. Already, limits on the GW background are starting to limit some inflation-era and cosmic string models. There seems little doubt that the proposed Square Kilometer Array radio telescope will be able to not only detect GW but to also study the properties of the GW sources in some detail, opening up a new era in astrophysics. \begin{theacknowledgments} The Parkes Pulsar Timing Array project has a great team of people based at the collaborating institutions. They contribute in different ways to helping us achieve our goals and I thank them all for their efforts. The Parkes radio telescope is part of the Australia Telescope which is funded by the Commonwealth of Australia for operation as a National Facility managed by CSIRO. \end{theacknowledgments} | 7 | 10 | 0710.5026 |
0710 | 0710.1480_arXiv.txt | {} {By reevaluating a 13-month stretch of Ulysses SWICS H pickup ion measurements near 5~AU close to the ecliptic right after the previous solar minimum, this paper presents a determination of the neutral interstellar H density at the solar wind termination shock and implications for the density and ionization degree of hydrogen in the LIC.} {The density of neutral interstellar hydrogen at the termination shock was determined from the local pickup ion production rate as obtained close to the cut-off in the distribution function at aphelion of Ulysses. As shown in an analytical treatment for the upwind axis and through kinetic modeling of the pickup ion production rate at the observer location, with variations in the ionization rate, radiation pressure, and the modeling of the particle behavior, this analysis turns out to be very robust against uncertainties in these parameters and the modeling. } {Analysis using current heliospheric parameters yields the H density at the termination shock equal to $0.087\pm0.022$~cm$^{-3}$, including observational and modeling uncertainties. } {} | Neutral interstellar gas of the local interstellar cloud (LIC) penetrates into the inner heliosphere as a neutral wind due to the relative motion between the Sun and the LIC. Apparently, the Sun is found near the boundary of a warm, relatively dilute cloud of interstellar gas, possibly with a significant gradient in the ionization fraction of H and He \citep[e.g.][]{cheng_bruhweiler:90a, wolff_etal:99, slavin_frisch:02} within a very structured surrounding \citep[e.g. reviews by][]{cox_reynolds:87a, frisch:95a}. In a companion paper within this special section, Frisch and Slavin (2008) lay out how the physical parameters and composition of the LIC at the location of the Sun, as derived from in-situ observations and from absorption line measurements, constrain the ionization state and radiation environment of the LIC. In situ observations of the two main constituents of the LIC, H and He, have been obtained with increasing accuracy, starting with the analysis of backscattered solar Lyman-$\alpha$ intensity sky maps \citep{bertaux_blamont:71, thomas_krassa:71} for H as well as with rocket-borne \citep{paresce_etal:74a} and satellite-borne \citep{weller_meier:74} observations of interstellar He using the solar He I 58.4 nm line. The optical diagnostics was followed by discovery of pickup ions for He \citep{mobius_etal:85a} and for H \citep{gloeckler_etal:92} and finally by direct neutral He observations \citep{witte_etal:93}. Such in-situ diagnostics, even at 1~AU, is made possible by the neutral gas flow deep into the inner heliosphere. Through the interplay between this wind, the ionization of the neutrals upon their approach to the Sun, and the Sun's gravitational field (distinctly modified by radiation pressure for H) a characteristic flow pattern and density structure is formed, with a cavity close to the Sun and gravitational focusing on the downwind side (for all species except H). The basic understanding of the related heliosphere -- LIC interaction has been summarized in early reviews by \citet{axford:72, fahr:74, holzer:77, thomas:78}. While He provides us with almost completely unbiased information about the physical parameters of the LIC since it enters the heliosphere unimpeded, the abundance of H and O, along with other species, is significantly depleted, their speed decreased, and their temperature increased through charge exchange in the heliospheric interface \citep{fahr:91, rucinski_etal:93a, izmodenov_etal:99a, mueller_etal:00, izmodenov_etal:04a}. A consolidation of the physical parameters of interstellar He, including the flow velocity vector relative to the Sun, as determined from neutral gas, pickup ion, and UV backscattering observations, was achieved through the effort of an ISSI Team \citep[][and references therein]{mobius_etal:04a}, thus leading to a benchmark for the physical parameters of the LIC. This paper is part of a follow-up effort within an ISSI Team to also consolidate the determination of the LIC H density. The determination of the H density in the LIC proper not only involves a measurement inside the heliosphere, but is also dependent on the filtration of H in the heliospheric boundary. Therefore, consolidating the observational results concentrates on the determination of the H density at the termination shock, which still requires taking into account of the dynamics of the flow into the inner heliosphere as well as ionization effects. This paper deals with the determination of the H density from the pickup ion observations made with Ulysses SWICS. The effort to obtain the density from mass-loading of the solar wind by H pickup ions and its resulting slowdown at large distances from the Sun is described in the paper by \citet[this volume]{richardson_etal:08a}, while \citet[this volume]{pryor_etal:08a} discuss a determination of the H density based on the reduction of the modulation of the UV backscatter signal with distance from the Sun. To illustrate the state of the model-dependent H density value in the LIC \citet[this volume]{mueller_etal:08a} compare different global models of the heliosphere and their results for the distances of the key boundary structures and the filtration factor, which also connects the inner heliosphere observations to the ionization state of the LIC, discussed by \citet[this volume]{slavin_frisch:08a}. In their previous work \citet{gloeckler_geiss:01b} used pickup ion fluxes as observed at 5~AU with Ulysses SWICS, the charge exchange rates from SWOOPS, and a Vasyliunas \& Siscoe distribution function to deduce the local neutral H density; they then used a hot interstellar gas model with the ionization rate significantly modified by electron ionization to deduce the density at the termination shock. In the present paper we use the same data set, but follow a complementary approach. After discussing previous derivations of the local neutral gas density and its extrapolation to the termination shock at the beginning of section 2 we present an alternative approach. We make use of the fact that the ionization rate in the pickup ion production appears both as production rate of PUI and as loss rate of the parent neutral gas population and any variations balance close to the aphelion of Ulysses. As a consequence, the PUI production rate is almost exactly proportional to the H density at the termination shock. In the same section we illustrate this behavior in a simplified analytical model that applies to the upwind region. In section 3 we simulate the local PUI production rate, starting with the density in the interstellar medium, compare it with the observations, and confirm the robustness of this approach by varying the parameters. We start with Monte-Carlo simulations of the flow through the heliospheric interface for two different LIC parameter sets, which result in two different H densities at the nose of the termination shock. In a second step, we hand these results over to a 3D time dependent test-particle code to calculate the H densities and H$^+$ PUI production rates at Ulysses during the observation interval, while accurately taking into account losses and radiation pressure along the trajectories of interstellar gas. We find the density at the termination shock and the LIC parameters that fit the observations best by interpolating between the two initial models. We study the response of the resulting PUI production rates to variations in the ionization rate, radiation pressure, and details in the modeling. In section 4, we present the results and show that our method is very robust against uncertainties of these parameters, and, in fact, against details of the simulations. \begin{figure} \resizebox{\hsize}{!}{\includegraphics[width=8cm]{aanda07a2_gr3.eps}} \caption{The interface correction: ratio of results of the Moscow Monte Carlo model for the geometry of Ulysses H PUI observations to the results of a sum of classical hot models (two populations), evaluated for identical parameters of the gas at the nose of the termination shock and identical ionization rate and radiation pressure as used in the MC model. The dotted horizontal line marks the heliospheric correction value equal to 1, the shaded area corresponds to the range of Ulysses heliocentric distances during the observations.} \label{aa} \end{figure} \begin{figure} \resizebox{\hsize}{!}{\includegraphics[width=8cm]{aanda07a3_gr1.eps}} \caption{Change of Ulysses position during the observations. The horizontal axis shows the heliocentric distance, the left-hand vertical scale heliolatitude, and the right-hand vertical scale ecliptic latitude. The times are indicated at the plot. Ecliptic longitude varied from $153.6\degr$ at the beginning of observations interval to $157.5\degr$ at the end, which corresponds to a change from $78.4\degr$ to $81.8\degr$ in heliolongitude. } \label{obs1} \end{figure} \begin{figure} \resizebox{\hsize}{!}{\includegraphics[width=8cm]{puiProdProf.eps}} \caption{Absolute production rates of H PUI as a function of heliocentric distances, obtained from the observed PUI spectrum during the interval discussed in the paper. Shown are the rates after normalizing the spectrum to the production rate at 1~AU and then multiplying by $1/r^2$. } \label{obs2} \end{figure} | We have used the accumulation of the H$^+$ pickup ion production rate from SWICS/Ulysses over a $\sim 13$~month period in 1997 -- 1998 at the Ulysses passage through the solar equator plane to infer the interstellar H density at the termination shock. By extensive simulations we demonstrated that the H$^+$ PUI production rate in this location of the heliosphere is only weakly dependent on the values of solar radiation pressure and neutral H ionization rate, but sensitively depend on the density at the termination shock. We have found that the H density at the termination shock inferred from these pickup ion production rates is very robust against any variations in the ionization rate, radiation pressure, and the actual modeling approaches for the density distribution in the inner heliosphere. In the present analysis we have, for the first time, included explicitly both the observational uncertainty and the modeling uncertainties. While in our approach the modeling uncertainties are minimized to a few percent, a larger uncertainty is incurred for the observation because absolute flux values are used, which results in an uncertainty of the obtained termination shock density equal to $\pm 25$\%. In the previous approach by \citet{gloeckler_geiss:01b} the observational uncertainty was minimized by making use of the ratio of the pickup ion and solar wind flux, but the $\sim 10$\% uncertainty quoted in \citet{izmodenov_etal:03a} does not include a range of values for the ionization rate and radiation pressure and the uncertainty of the geometric factor of the instrument. But, as demonstrated in the previous sections, the density at the termination shock scales linearly with the observed pickup ion production rate, which is directly related to the pickup ion flux and/or distribution function. Hence, any observational uncertainties will transfer linearly into the resulting densities. Since \citet{izmodenov_etal:03a} started from the local neutral gas density at Ulysses, any uncertainty in the ionization rate will appear approximately linearly in the extrapolated density at the termination shock. The determination of the PUI production rate at Ulysses near the aphelion, on which our derivation of $n_{\mathrm{H,TS}}$ is based, is not entirely model-free. Although in the present determination of the PUI production rate at Ulysses a simple hot model was used for forward-modeling of the PUI distribution function, our method is robust against simplifications inherent to that kind of modeling because it uses a quantity (i.e., the production rate of PUI at Ulysses) which weakly depends on details of such modeling. The determination of the H density at the termination shock by \citet{gloeckler_etal:08a}, equal to $0.080\pm0.008$~cm$^{-3}$, is free of the uncertainty of the geometric factor of the instrument, but is subject to a combination of uncertainties in the determination of the He density and that of the He abundance relative to H from the Voyager LECP observations, including uncertainties in the ratios of the production rates of these species. Nevertheless, after including of all uncertainties, all three approaches (i.e. the present one and those from \citet{gloeckler_geiss:01b} and \citet{gloeckler_etal:08a}) should be read with a similar uncertainty band. The density value presented here agrees very well with the new determination by \citet{gloeckler_etal:08a}. Although our value still agrees with the previous determination of the H density from SWICS pickup ion observations by \citet{gloeckler_geiss:01b} within their mutual uncertainty bands, the combination of the two new results suggest a somewhat lower density than 0.1~cm$^{-3}$. Our results also agree comfortably with the TS density values found from the analysis of the heliospheric Lyman-$\alpha$ glow \citep[this volume]{pryor_etal:08a} and from the solar wind slowdown \citep[this volume]{richardson_etal:08a} within the uncertainty bands. It should be noted here that the coupling between the neutral and ionized component of the interstellar medium between the bow shock and the heliopause appears to be somewhat stronger than suggested previously. This is a consequence of an updated relation for the energy dependence of the charge exchange cross-section between protons and H atoms \citep{lindsay_stebbings:05a}. The parameters of the interstellar gas in front of the heliospheric bow shock, as assumed in simulation (1), also seem to be robust, as shown by \citet[this volume]{mueller_etal:08a}, who discussed the present status of the modeling of heliospheric interface and showed that differences in the filtration rate returned by different models of the heliosphere evaluated with identical initial parameters are about 15\% and this result can be adopted as the uncertainty of the H density in the CHISM. \begin{appendix} | 7 | 10 | 0710.1480 |
0710 | 0710.1952_arXiv.txt | We present a Markov Chain Monte Carlo global analysis of neutrino parameters using both cosmological and experimental data. Results are presented for the combination of all presently available data from oscillation experiments, cosmology, and neutrinoless double beta decay. In addition we explicitly study the interplay between cosmological, tritium decay and neutrinoless double beta decay data in determining the neutrino mass parameters. We furthermore discuss how the inference of non-neutrino cosmological parameters can benefit from future neutrino mass experiments such as the KATRIN tritium decay experiment or neutrinoless double beta decay experiments. | The question of neutrino mass is one of the most profound in modern particle physics. Most plausible models of neutrino mass solve the puzzle of why neutrino masses are so small by introducing a new scale at high energy, and precision studies of neutrino physics therefore hold the potential to investigate physics at scales beyond those reachable in current accelerator experiments. They also make the study of the possible Majorana nature of neutrinos possible (see \cite{Mohapatra:2004ht,Mohapatra:2004vr,Mohapatra:2006gs} for a thorough discussion of this). While the neutrino mass differences have now been measured at about 10\% precision by oscillation experiments (see e.g.\ \cite{Maltoni:2004ei,fogli}) the absolute mass scale remains unknown and inaccessible to oscillation experiments. There are, however, several possible paths to measuring the absolute neutrino mass. The kinematical effect of neutrino mass can be probed either via its effect on the beta decay spectrum or via its effect on cosmological structure formation. If neutrinos are Majorana particles a different possibility is to search for neutrinoless double beta decay because the transition probability for this process is proportional to the neutrino mass squared. In the past year there have been several papers discussing how to unify the data analysis for the various approaches \cite{fogli,Host:2007wh}. This is a non-trivial issue, given that completely different physics is involved and that the three probes are actually sensitive to three distinct observables. Here we present a new Markov Chain Monte Carlo global analysis of neutrino parameters using both cosmological and experimental data. The analysis software is based on the CosmoMC Markov Chain Monte Carlo (MCMC) package for cosmological parameter estimation \cite{Lewis:2002ah,cosmomc}, appropriately modified to incorporate all parameters related to neutrino physics. This approach uses Bayesian inference instead of the frequentist method commonly used in particle physics. The approach is somewhat similar to the MCMC technique developed in \cite{trotta} to constrain MSSM parameters. However, a key difference is that here we keep the full cosmological parameter estimation which allows for a closer study of the interplay between neutrino data and cosmological parameter estimation. In Section II we describe the methodology used and in Section III we present the main results for various different assumptions about present and future data, as well as different parameter spaces. Finally we present our conclusions in Section IV. | A detailed neutrino parameter estimation study has been carried out using the Markov Chain Monte Carlo technique with the goal of unifying the various techniques for measuring the absolute neutrino mass scale. The MCMC technique is extremely powerful in this regard and allows for a very fast scanning many-dimensional likelihood spaces. In the concrete example here we have used 8 parameters describing the properties of light, active Majorana neutrinos, and 6 further parameters which specify the cosmology. We find that for present data the combination of cosmological data with the upper limit on $m_{\beta \beta}$ from Heidelberg-Moscow slightly improves the existing cosmological bound on the sum of neutrino masses. More interestingly we have studied the interplay between various future constraints from cosmology, tritium decay and neutrinoless double beta decay. If all probes come up with a negative result the addition of data sets does not yield any radically new information. However, we have also studied an example in which the upcoming KATRIN and GERDA experiments are both assumed to provide tentative evidence for neutrino mass. In this case the combination of all three types of data allows for a much stronger constraint on neutrino properties than otherwise allowed. Finally we have also studied how experimental data from tritium decay or neutrinoless double beta decay can help in cosmological parameter estimation, particularly concerning the dark energy equation of state. It should be noted that in the present analysis only presently available cosmological data has been used. In the same time frame as KATRIN and GERDA new cosmological data will become available and is likely to improve the cosmological neutrino mass bound significantly (see \cite{Xia:2007gz,Gratton:2007tb,Hannestad:2007cp,Perotto:2006rj,Takada:2005si,% Lesgourgues:2005yv,Wang:2005vr,Song:2003gg,Hannestad:2002cn} for a non-exhaustive list). In the somewhat longer term cosmological constraints can be potentially be pushed below 0.1 eV sensitivity to $\sum m_\nu$. At the same time neutrinoless double beta decay experiments will have equally improved sensitivity and it will very likely be possible to determine the absolute neutrino mass as well as the nature of the mass hierarchy. In conclusion, the combination of cosmological data with experimental neutrino data in a global analysis will be extremely useful in the future, when more precise experimental data becomes available. | 7 | 10 | 0710.1952 |
0710 | 0710.3166_arXiv.txt | Despite intense scrutiny, the progenitor system(s) that gives rise to Type Ia supernovae remains unknown. The favored theory invokes a carbon-oxygen white dwarf accreting hydrogen-rich material from a close companion until a thermonuclear runaway ensues that incinerates the white dwarf. However, simulations resulting from this single-degenerate, binary channel demand the presence of low-velocity H$\alpha$ emission in spectra taken during the late nebular phase, since a portion of the companion's envelope becomes entrained in the ejecta. This hydrogen has never been detected, but has only rarely been sought. Here we present results from a campaign to obtain deep, nebular-phase spectroscopy of nearby Type Ia supernovae, and include multi-epoch observations of two events: SN~2005am (slightly subluminous) and SN~2005cf (normally bright). No H$\alpha$ emission is detected in the spectra of either object. An upper limit of $0.01\ M_\odot$ of solar abundance material in the ejecta is established from the models of \citet{Mattila05} which, when coupled with the mass-stripping simulations of \citet{Marietta00} and \citet{Meng07}, effectively rules out progenitor systems for these supernovae with secondaries close enough to the white dwarf to be experiencing Roche lobe overflow at the time of explosion. Alternative explanations for the absence of H$\alpha$ emission, along with suggestions for future investigations necessary to confidently exclude them as possibilities, are critically evaluated. | \label{sec:1} Ever since the currently favored single-degenerate, binary channel was proposed as the progenitor system for Type Ia supernovae \citep{Whelan73}, all models of the impact of the exploded white dwarf (WD) on the secondary star have indicated that significant amounts of solar-abundance material, stripped from the secondary's envelope, become entrained in the ejecta \citep{Wheeler75,Fryxell81,Taam84,Chugai86,Livne92,Marietta00,Meng07}. Observational evidence for this material, however, still eludes us \citep{Mattila05}, and serves as one reminder among many that we still lack direct observational proof for the single-degenerate scenario \citep{Branch95,Livio01}. The most detailed theoretical investigation of the expected amount and distribution of stripped material within a young Type Ia supernova (SN~Ia) remnant is that of \citet{Marietta00}, who studied the problem with two-dimensional numerical simulations. Four basic progenitor systems were investigated, including three with secondaries (a main-sequence, subgiant, or red giant star) close enough for mass-transfer to occur through Roche lobe overflow (RLOF), and one containing a secondary (a red giant) donating material through a strong stellar wind --- the symbiotic case. All secondaries were given masses of $\sim 1\ M_\odot$ at the time of the explosion and were placed either just within (the RLOF cases), or just beyond (the symbiotic case), the limiting distance from the WD within which RLOF can occur (i.e., $a/R = 3$, where $a$ is the orbital separation in units of the secondary star's radius, $R$; see \citealt{Eggleton83}). Three additional systems, in which a main-sequence secondary was placed too far away to experience RLOF (thus rendering it an unlikely progenitor system) were also included to establish the scaling between orbital separation and amount of stripped material. The numerical results of the \citet{Marietta00} study confirmed predictions from earlier analytic work \citep[e.g.,][]{Wheeler75,Chugai86} that substantial material is indeed stripped, with the amount ranging from a minimum of $0.15\ M_\odot$ for a close ($a/R = 3$) main-sequence secondary, up to nearly the entire envelope ($\sim 0.5\ M_\odot$) for a similarly placed red giant. Increasing the orbital separation beyond the RLOF limit in the \citet{Marietta00} simulations resulted in a dramatic decrease in the amount of stripped material: For $a/R = 12$, a main-sequence secondary loses only $0.0018\ M_\odot$. However, since such systems lack an efficient mechanism for mass transfer, they are not considered to be viable SN~Ia progenitors. A red giant secondary placed at a similarly distant location, however, could potentially donate mass through a strong stellar wind, making the symbiotic case a possibility if large orbital separations are required \citep[e.g.,][]{Munari92}. As discussed by \citet{Meng07}, one shortcoming of the \citet{Marietta00} study is their use of standard solar-model stars for the companion, rather than companions whose structures have been appropriately modified due to having evolved in a binary system \citep[e.g.,][]{Eggleton73}. To investigate the effect this simplification might have had on the results, \citet{Meng07} use an analytic model to estimate the amount of mass expected to be stripped from a variety of evolved secondaries. (Their analytic approach was first tested using the unevolved secondaries used by \citealt{Marietta00}, and was demonstrated to approximate the results obtained numerically.) The result is that the quantity of material expected to be stripped from evolved secondaries is considerably lower than that predicted for standard solar-model companions. In fact, \citet{Meng07} find that the minimum value for systems experiencing RLOF is diminished from $0.15\ M_\odot$ to only $0.035\ M_\odot$. The reduction arises primarily from the pre-explosion mass-loss producing a more compact companion star whose material is more difficult to strip than it is in the unevolved case. \citet{Meng07} stress, however, that their new values are really lower bounds on the amount of stripped material, since their analytic approach does not consider the thermal energy imparted by the ejecta to the companion envelope, which likely serves to heat and vaporize a portion of it and thereby increase the amount of stripped material \citep[e.g.,][]{Fryxell81,Mattila05}. Thus, $0.035\ M_\odot$ serves as a conservative lower bound on the expected amount of stripped material resulting from their models. The typical velocity of the stripped material is found in all studies to be far slower than the $\sim 10,000$~\kms\ velocity that characterizes the bulk of the iron-rich ejecta \citep{Chugai86,Marietta00,Meng07}. This has the effect of placing it almost entirely in the central region of the supernova remnant, with the majority of it predicted to be moving with a velocity of under $1,000$~\kms\ \citep{Marietta00}. The stripped material is largely confined to the downstream region behind the companion star, where it contaminates a solid angle that ranges from $66^\circ$ for the main-sequence companion to $115^\circ$ for the red-giant companion \citep{Marietta00}. The low velocity of the stripped material renders it undetectable when the faster-moving, iron-rich ejecta are optically thick. However, detailed radiative transfer calculations performed by \cite{Mattila05} predict that it should become visible at late times, when the outer ejecta have thinned out and become transparent enough to reveal the slower-moving gas in the central regions. The most prominent expected spectral signature of the companion star's stripped material is narrow H$\alpha$ emission in nebular spectra \citep{Mattila05}, taken more than $\sim 250$ days after maximum light. The H$\alpha$ emission should be present within $\pm 1,000$~\kms\ {\rm of\ } $\lambda_0 = 6563$ \AA\ but, due to the expected asymmetry in the distribution of the solar-abundance material, could present an H$\alpha$ profile ranging from a very narrow spike to a broader emission line in the observed spectrum. Obtaining spectra with high enough resolution to resolve fairly narrow lines is thus a useful component of a targeted search for this H$\alpha$. To date, only a few nebular SN~Ia spectra have been obtained and H$\alpha$ has never been detected, although the majority of the spectra lack the spectral resolution and signal-to-noise ratio to place interesting constraints on the companion. The best limits, by far, come from the recent study by \cite{Mattila05}, which constrains hydrogen-rich material in SN~2001el to be $\lesssi 0.03\ M_\odot$ from a low-resolution ($\sim 700$~\kms) spectrum obtained 398 days past maximum light. In an effort to expand and improve on earlier work, both in terms of the number of objects studied as well as the resolution, sensitivity, and temporal coverage of the spectra, we have initiated a program to obtain deep, moderate resolution ($\lesssim 150$~\kms, or $\sim 3$~\AA, at H$\alpha$), late-time spectra of SNe~Ia at multiple epochs using the Keck and Gemini telescopes. The first phase of this project has garnered data on two objects: SN~2005am, a slightly subluminous \citep{Li06} event and SN~2005cf, an SN~Ia of normal brightness \citep{Pastorello07}. We present and analyze our observations in \S~\ref{sec:2}, discuss the results in \S~\ref{sec:3}, and conclude in \S~\ref{sec:4}. | \label{sec:4} We obtained five deep, moderate-resolution, nebular-phase spectra of two SNe~Ia (SN~2005am and SN~2005cf) in order to search for narrow H$\alpha$ emission that would betray the existence of material stripped from the envelope of a mass-donating stellar companion to the exploding WD. No such emission is detected in either object at any epoch. From the models of \citet{Mattila05}, we establish upper limits of $0.01\ M_\odot$ of solar abundance material in the inner ejecta of both objects, which are the tightest constraints yet established by such studies. Our non-detections of H$\alpha$, coupled with the mass-stripping results of \citet{Marietta00} and \citet{Meng07}, rule out all hydrogen-donating companions close enough to the WD to have been experiencing RLOF at the time of explosion for these events. Additional theoretical work is needed in several areas to buttress this conclusion, including most critically verification of the transparency of the outer, more rapidly moving ejecta that could potentially block H$\alpha$ photons from escaping from the inner region. Bearing this caveat in mind, we propose that symbiotics are, at this time, the most likely progenitor class that remains consistent with these data. Definitive proof of the identity of the progenitor system(s) that gives rise to SNe~Ia remains elusive, and it must be admitted that our conclusion, which is based on the {\it lack} of a detection, is not as satisfying as one based {\it on} a detection. Should future modeling efforts prove unable to ``hide the hydrogen'' for even widely separated binaries, then the continued viability of the single-degenerate, hydrogen-donating progenitor will require that H$\alpha$, no matter how weak, must ultimately be detected. | 7 | 10 | 0710.3166 |
0710 | 0710.3399_arXiv.txt | There is currently no explanation of why the corona has the temperature and density it has. We present a model which explains how the dynamics of magnetic reconnection regulates the conditions in the corona. A bifurcation in magnetic reconnection at a critical state enforces an upper bound on the coronal temperature for a given density. We present observational evidence from 107 flares in 37 sun-like stars that stellar coronae are near this critical state. The model may be important to self-organized criticality models of the solar corona. | Dynamics in the solar corona takes on a wide range of forms. On one hand, the corona is the setting for the most violent eruptions in the solar system: solar flares and coronal mass ejections \citep{Aschwanden01}. On the other, coronal heating makes the corona almost a thousand times hotter than the photosphere, even in the quiet sun \citep{Klimchuk06}. \citet{Parker83,Parker88} unified these two phenomena by proposing that micro- and nano-flares, less energetic cousins of eruptive flares, heat the corona. This model gained credence from studies showing that solar flares exhibit power law statistics \citep{Lin84,Dennis85,Crosby93,Feldman97,Wheatland00, Nita02,Paczuski05} over a wide range of scales for many quantities. [See \citet{Charbonneau01} for a review.] In addition, stellar flares have similar light curves to solar flares \citep{Gershberg05} and also exhibit power law statistics \citep{Collura88,Shakhovskaya89, Audard00}, suggesting that the physics of the solar corona is generic to sun-like stars. Coronal dynamics remains an active research area \citep{Hudson91, Georgoulis98,Shibata02,Hughes03}. Details of the eruption process including how magnetic energy is stored, how eruptions onset, and how the stored energy is converted to other forms are still open questions. In addition, while micro- and nano-flares are believed to be a major contributor to coronal heating, the authors know of no theory which explains why the coronal temperature and density have the values they have, as opposed to larger or smaller values. In this paper, we propose that the condition of the corona is regulated by magnetic reconnection \citep{Cassak06b}, a dynamical process which converts magnetic energy into kinetic energy and heat and energizes particles. Magnetic energy is stored during collisional (slow) reconnection, which has been shown to drive the coronal plasma toward lower collisionality \citep{Cassak06}. If the plasma becomes marginally collisionless, a bifurcation in the underlying dynamics of reconnection occurs \citep{Cassak07c}. This bifurcation, which occurs when two length scales $\delta_{SP}$ and $\rho_{i}$ (to be defined below) are comparable, catastrophically initiates fast (Hall) reconnection, releasing the stored energy in the form of an eruption. The condition of marginal collisionality, therefore, sets an upper bound on how hot the coronal plasma can be for a given density. The continual driving toward lower collisionality of the pre-flare corona by slow reconnection enforces the self-organization of the corona to a state of marginal collisionality where $\delta_{SP} \sim \rho_{i}$. We present finer details of this process below. Then, we perform the first observational test of this model using a large sample of data from stellar flares on sun-like stars. We find that $\delta_{SP}$ and $\rho_{i}$ are comparable for every event in the sample, indicating that stellar coronae do self-organize into a marginally collisional state. | \label{sec-disc} The data analyzed in this paper pertain to flares in sun-like stars, but the underlying dynamics of reconnection is general. Our model applies equally well to micro- and nano-flares in the quiet corona. Using values for the quiet sun of $T \sim$ 1 MK, $n \sim 10^{9}$ ${\rm cm}^{-3}$, $B \sim $ 5 G, and $L \sim 10^{10}$ cm, we find $\delta_{SP} \sim 770$ cm and $d_{i} \sim 720$ cm, in agreement with the model. The present result may have important implications for self-organized criticality (SOC) models of the solar corona. SOC occurs in driven, dissipative systems when the system is driven to a critical state where it undergoes a major reconfiguration \citep{Bak87}. SOC leads to power law statistics, which encouraged \citet{Lu91} to propose the corona undergoes SOC. Subsequent studies of SOC in the corona exist \citep{Lu93,Vlahos95,Longcope00,Isliker01}, but a firm physical foundation of the mechanism for self-driving and the physical condition setting the critical state is often traded for the ease of performing cellular automaton simulations [see \citet{Charbonneau01} for a review]. The present result provides a physical mechanism for self-driving (embedded Sweet-Parker reconnection) and the critical state (marginal collisionality), which may provide an avenue for developing quantitative predictions of SOC to compare with coronal observations. An alternate mechanism \citep{Uzdensky06,Uzdensky07,Uzdensky07b} for heating the solar corona uses a change in density to achieve self-regulation. After an eruption, chromospheric evaporation increases the coronal density, decreasing the ion gyroradius [eq.~(\ref{didef})] and making subsequent eruptions more difficult. The extent to which Uzdensky's and our mechanisms regulate coronal heating is an open question. The present model assumes that Sweet-Parker scaling is appropriate for thin current sheets of large extent. Long current sheets are known to fragment due to secondary instabilities, but the effect of this on the reconnection rate is unknown. Verification of the present model would entail testing whether Sweet-Parker reconnection in extended current sheets remains much slower than Hall reconnection. [See \citet{Uzdensky07b} for further discussion of this point as well as other future research directions.] The authors thank J.~F.~Drake, A.~Klimas, E.~Ott, S.~Owocki, P.~So and D.~Uzdensky for helpful conversations. This work was supported in part by the Delaware Space Grant. \clearpage | 7 | 10 | 0710.3399 |
0710 | 0710.5789_arXiv.txt | Recent stellar evolutionary calculations of low-metallicity massive fast-rotating main-sequence stars yield iron cores at collapse endowed with high angular momentum. It is thought that high angular momentum and black hole formation are critical ingredients of the collapsar model of long-soft $\gamma$-ray bursts (GRBs). Here, we present 2D multi-group, flux-limited-diffusion MHD simulations of the collapse, bounce, and immediate post-bounce phases of a 35-\msun\, collapsar-candidate model of Woosley \& Heger. We find that, provided the magneto-rotational instability (MRI) operates in the differentially-rotating surface layers of the millisecond-period neutron star, a magnetically-driven explosion ensues during the proto-neutron star phase, in the form of a baryon-loaded non-relativistic jet, and that a black hole, central to the collapsar model, does not form. Paradoxically, and although much uncertainty surrounds stellar mass loss, angular momentum transport, magnetic fields, and the MRI, current models of chemically homogeneous evolution at low metallicity yield massive stars with iron cores that may have {\it too much} angular momentum to avoid a magnetically-driven, hypernova-like, explosion in the immediate post-bounce phase. We surmise that fast rotation in the iron core may inhibit, rather than enable, collapsar formation, which requires a large angular momentum not in the core but {\it above} it. Variations in the angular momentum distribution of massive stars at core collapse might explain both the diversity of Type Ic supernovae/hypernovae and their possible association with a GRB. A corollary might be that, rather than the progenitor mass, the angular momentum distribution, through its effect on magnetic field amplification, distinguishes these outcomes. | There is mounting observational evidence for the association between long-soft $\gamma$-ray Bursts (GRBs) and broad-lined Type Ic supernovae (SNe; see Woosley \& Bloom 2006 for a review). Such hydrogen-deficient (and, perhaps, also helium-deficient) progenitors are compact and, if fast rotating in their % core at collapse, fulfill critical requirements for the formation of a collapsar (Woosley 1993). The engine that converts energy from long-term accretion of disk material onto the black-hole (BH) may power a relativistic jet in the excavated polar regions. The jet breaks out of the progenitor surface while equatorial accretion continues. Depending on the BH mass and the angular momentum budget in the collapsing envelope, this ``engine'' may operate for seconds, i.e. as long as typical long-soft GRBs. Accompanying this beamed relativistic polar jet might be a disk wind, fueled by neutrinos or MHD processes, which would explode the Wolf-Rayet envelope. This explosion and the radioactive $^{56}$Ni material produced might lead to very energetic, broad-lined, Type Ic SN of the hypernova variety (Iwamoto et al. 1998; MacFadyen \& Woosley 1999; Hjorth et al. 2003; Stanek et al. 2003). State-of-the-art radiation-hydrodynamics simulations including a sophisticated equation of state (EOS) and detailed neutrino transport (Buras et al. 2003; Burrows et al. 2006,2007a; Kitaura et al. 2006; Marek \& Janka 2007; Mezzacappa et al. 2007) suggest that while the neutrino mechanism of supernova explosions may work for the lower-mass massive progenitors, it may not for the more massive progenitors, characterized by an ever higher post-bounce accretion rate onto the proto-neutron star (PNS). Burrows et al. (2006,2007a) have suggested that an acoustic mechanism will work for all slowly rotating progenitors that do not explode by other means within the first second after bounce. However, massive star cores endowed with a large angular momentum at the time of collapse should experience the magneto-rotational instability (MRI; Balbus \& Hawley 1991; Akiyama et al. 2003; Pessah et al. 2006; Shibata et al. 2006; Etienne et al. 2006), with the potential to exponentially amplify weak initial fields on a rotation timescale. The saturation values of such fields are ultimately set by the free-energy of differential rotation available in the surface layers of the PNS (Ott et al. 2006), and can be large, i.e., on the order of 10$^{15}$\,G at a radius of a few tens of kilometers. The corresponding magnetic stresses at the neutron star surface lead systematically to powerful jet-like explosions $\sim$100\,ms after bounce (see, e.g., Ardeljan et al. 2005; Yamada \& Sawai 2004; Kotake et al. 2004; Sawai et al. 2005; Moiseenko et al. 2006; Obergaulinger et al. 2006; Burrows et al. 2007b, hereafter B07; Dessart et al. 2007). In this letter, we investigate, in the context of the collapsar model, the potential implications of this magnetic explosion mechanism. Our study focuses on the immediate post-bounce phase, whose importance was emphasized by Wheeler et al. (2000,2002). This is in contrast to previous work which explored only the phase subsequent to BH formation (MacFadyen \& Woosley 1999; Aloy et al. 2001; Zhang et al. 2003; Proga 2005). Indeed, two terms sometimes used in the collapsar context are ``failed SN'' (MacFadyen \& Woosley 1999) and ``prompt BH formation'' (MacFadyen et al. 2001). Our analysis supports the idea that the conditions for the collapsar model, as stated so far, are also suitable for a magnetically-driven explosion in the immediate post-core-bounce PNS phase, and that BH formation may be so delayed for a range of putative progenitor models that it does not in fact occur\footnote{In the present context, BH formation is never prompt, since it takes a finite time, on the order of seconds, for the PNS to accumulate the critical mass at which it experiences the gravitational instability. This is in contrast with super-massive stars, such as the progenitors of pair-instability SNe, which may form an apparent horizon during collapse and thus ``directly'' transition to a BH (Liu et al. 2007).}. In \S2, we present radiation MHD simulations with the code VULCAN/2D (Livne et al. 2004,2007) of a collapsar-candidate model that support this thesis. In \S3, we discuss the implications of our results for stellar evolutionary models that might lead to collapsars and/or hypernovae. \begin{deluxetable}{lcccccccc} \tablewidth{0pt} \tabletypesize{\scriptsize} \tablecaption{Properties of our two MHD-VULCAN/2D simulations of the 35OC collapsar model of WH06.\label{tab_model}} \tablehead{ \colhead{}& \colhead{$t_{\rm end}$}& \colhead{$t_0$}& \colhead{M$_{10}$}& \colhead{P$_{10}$}& \colhead{E$_{\rm expl}$}& \colhead{\edot$_{\rm gas}$}& \colhead{\edot$_{\rm \vec{E} \times \vec{B}}$}& \colhead{$v_{\rm max}$}\\ \colhead{}& \colhead{ms}& \colhead{ms}& \colhead{\mo}& \colhead{ms}& \colhead{B}& \multicolumn{2}{c}{B\,s$^{-1}$}& \colhead{km\,s$^{-1}$} } \startdata M0 & 369 & ... & 2.1 & 4 & 0.03 & 0.5 & 0.25 & 43,000 \\ M1 & 666 & 349 & 1.7 & 12 & 3.31 & 9.4 & 3.0 & 58,000 \\ \enddata \tablecomments{ $t_{\rm end}$ gives the time at the end of each simulation, while $t_0$ is the time when the rate of polar mass ejection first overcomes equatorial mass accretion. All quoted quantities in the table correspond to the final time in each simulation, while times are given with respect to core bounce. M$_{10}$ (P$_{10}$) corresponds to the total baryonic mass (average rotation period) inside the 10$^{10}$\,g\,cm$^{-3}$ isodensity contour. \edot$_{\rm gas}$ (\edot$_{\rm \vec{E} \times \vec{B}}$) is the Bernoulli (Poynting) power in the ejecta, obtained by integrating the corresponding flux over a shell with a radius of 500\,km. [See text for additional information.] } \end{deluxetable} | The potential for exponential growth on a rotational timescale of initial seed magnetic fields by the MRI (Shibata et al. 2006; Etienne et al. 2006), fueled by the free energy of core rotation, makes the initial angular momentum budget of the progenitor star the key parameter in determining the outcome during the immediate post-bounce phase (B07). A magnetically-driven, baryon-loaded, and non-relativistic, explosion is obtained for WH06's 35OC collapsar candidate model, evolved at low metallicity from a 35\,\msun\, fast-rotating main-sequence star. The explosion occurs $\sim$200\,ms after bounce and reaches $\sim$3\,B $\sim$400\,ms later. After an initial accretion phase, the steadily decreasing PNS mass reaches only $\sim$1.7\,\msun\, at the end of the simulation, and, thus, the quasi-steady explosion we observe suggests that BH formation is unlikely to occur. Moreover, baryon contamination prevents the ejecta from becoming relativistic. Note that the production of a GRB in the collapsar context is contingent on the gravitational collapse of the PNS to a BH. The recent stellar evolutionary calculations of Yoon \& Langer (2005), WH06, and Meynet \& Maeder (2007) of fast-rotating main-sequence objects at low metallicity systematically predict such fast-rotating cores at collapse. Starting from similar conditions for a 35-40\,\msun\, star, but using different mass-loss ``recipes,'' they obtain very similar rotational profiles in the inner core. % Allowing for anisotropic mass loss (Meynet \& Maeder 2007), a model of C. Georgy (2007, priv. comm.) suggests an even larger (by a factor of two) specific angular momentum in the inner 3\,\msun\, at the end of silicon core burning. % Despite the agreement between these different evolutionary computations, the magnetically-driven explosion and the ``failed'' BH formation described here are conditional on the uncertain treatment of mass loss, angular momentum transport, and magnetic processes (Spruit 2002) during the pre-collapse evolution. At very low metallicities, radiatively-driven winds of massive stars are inhibited by the lack of metals (Kudritzki 2002; Vink et al. 2001; Vink \& de Koter 2005), whose optically-thick lines intercept radiation momentum (Castor et al. 1975). Recent revisions downward of mass-loss rates due to clumping (Owocki et al. 1988; Bouret et al. 2005; Fullerton et al. 2006) suggest, however, the potential importance of episodic outbursts, akin to the 1843 giant eruption of Eta Carina (Smith \& Owocki 2006). The metallicity dependence of such phenomena is entirely unknown, mostly because the fundamental cause of the outburst remains a mystery. While line driving seems excluded, continuum driving of a porous medium at super-Eddington luminosities has been proposed by Owocki et al. (2004) as an alternative. Finally, mass loss in fast-rotating, and sometimes critically-rotating (Townsend et al. 2004), envelopes is complicated by the effects associated with centrifugal support, surface oblateness, and gravity-darkening (Cranmer \& Owocki 1995; Owocki et al. 1996), so that the mass-loss ``recipes'' used in stellar evolutionary models are not always substantiated by observational and theoretical evidence. At present, and in light of our simulations, it appears that chemically-homogeneous evolution of fast-rotating main-sequence massive stars at low metallicity systematically yields iron cores at collapse that may have {\it too much} angular momentum, a property that prevents the formation of a collapsar. Uncertainties in the modeling of the pre-collapse evolution may result, however, in slower-rotating iron cores\footnote{Note that the rotational energy $E_{\rm rot}$ is a stiff function of angular velocity $w$, i.e., $E_{\rm rot} \propto w^2$.} and, thus, might inhibit an early magnetically-driven explosion in favor of black hole, and perhaps collapsar, formation. \begin{figure} \includegraphics[width=8cm]{f2.ps} \caption{Colormap of the entropy at 666\,ms after bounce for model M1, overplotted with white iso-density contours (every decade downward from 10$^{10}$\,g\,cm$^{-3}$) and velocity vectors (length saturated to 15\% of the width of the display and corresponding to a velocity of 30,000\,km\,s$^{-1}$.) } \label{fig_still} \end{figure} We conclude that variations in the angular momentum distribution of pre-collapse massive stars may lead to different post-bounce scenarios. Non- or slowly-rotating progenitors may explode with weak/moderate energy ($\sles$1\,B) through a neutrino or an acoustic mechanism $\sles$1\,s after bounce, or may collapse to a BH. Objects with large angular momentum in the envelope, but little in the core, may proceed through the PNS phase, transition to a BH and form a collapsar with a GRB signature. Owing to the modest magnetic-field amplification above the PNS, a weak precursor polar jet may be launched, soon overtaken by a baryon-free, collimated relativistic jet. At the same time, the progenitor envelope is exploded by a disk wind, resulting in a hypernova-like SN with a large luminosity (large $^{56}$Ni mass). Finally, and this is what we conclude here, objects with large angular momentum in the core may not transition to a BH. Instead, and fueled by core-rotation energy, a magnetically-driven baryon-loaded non-relativistic jet is obtained without any GRB signature. The explosion has the potential of reaching energies of a few B to 10\,B, and for viewers along the poles of looking like a Type Ic hypernova-like SN with broad lines. For a viewer at lower latitudes, the delayed and less energetic explosion nearer the equator may look more like a standard Type Ic SN (H\"{o}flich et al. 1999). This volume-restricted jet-like explosion is dimmer, as the amount of processed $^{56}$Ni may be significantly less than the $\sim$0.5\,\msun\, obtained in the collapsar context (MacFadyen \& Woosley 1999). Hence, magnetic processes during the post-bounce phase of fast-rotating iron cores offer a potential alternative to collapsar formation and long-soft GRBs by producing non-relativistic non-Poynting-flux-dominated baryon-loaded hypernova-like explosions without any GRB signature. Importantly, while our study narrows the range over which the collapsar model may exist, it also offers additional routes to explain the existence of GRB/SN-hypernova events like SN 1998bw (Woosley et al. 1999), and hypernova events like SN 2002ap without a GRB signature (Mazzali et al. 2002). More generally, magnetic effects should naturally arise in the context of gravitational collapse and fast rotation. The resulting angular momentum of newly-formed BHs and magnetars, for example, would be reduced, perhaps considerably, by any prior magnetically-driven explosion, and, thus, may decrease the power of subsequent mass ejections from compact objects (see, e.g., Thompson et al 2004). | 7 | 10 | 0710.5789 |
0710 | 0710.5740_arXiv.txt | {} {The properties of the very high energy (VHE; E$>$100 GeV) $\gamma$-ray emission from the high-frequency peaked BL\,Lac PG\,1553+113 are investigated. An attempt is made to measure the currently unknown redshift of this object.} {VHE Observations of PG\,1553+113 were made with the High Energy Stereoscopic System (HESS) in 2005 and 2006. H+K (1.45$-$2.45$\mu$m) spectroscopy of PG\,1553+113 was performed in March 2006 with SINFONI, an integral field spectrometer of the ESO Very Large Telescope (VLT) in Chile.} {A VHE signal, $\sim$10 standard deviations, is detected by HESS during the 2 years of observations (24.8 hours live time). The integral flux above 300 GeV is $(4.6\pm0.6_{\rm stat}\pm0.9_{\rm syst}) \times 10^{-12}$ cm$^{-2}$\,s$^{-1}$, corresponding to $\sim$3.4\% of the flux from the Crab Nebula above the same threshold. The time-averaged energy spectrum is measured from 225 GeV to $\sim$1.3 TeV, and is characterized by a very soft power law (photon index of $\Gamma = 4.5\pm0.3_{\rm stat}\pm0.1_{\rm syst}$). No evidence for any flux or spectral variations is found on any sampled time scale within the VHE data. The redshift of PG\,1553+113 could not be determined. Indeed, even though the measured SINFONI spectrum is the most sensitive ever reported for this object at near infrared wavelengths, and the sensitivity is comparable to the best spectroscopy at other wavelengths, no absorption or emission lines were found in the H+K spectrum presented here.} {} | Evidence for VHE ($>$100 GeV) $\gamma$-ray emission from the active galactic nucleus (AGN) PG\,1553+113 was first reported by the HESS collaboration (\cite{HESS_discovery}) based on observations made in 2005. This evidence was later confirmed (\cite{MAGIC_1553}) with MAGIC observations in 2005 and 2006. Similar to essentially all AGN detected at VHE energies, PG\,1553+113 is classified as a high-frequency peaked BL\,Lac \cite{classification} and is therefore believed to possess the double-humped broad-band spectral energy distribution (SED) typical of blazars. The low-energy (i.e. from the radio to the X-ray regime) portion of the SED of PG\,1553+113 is well-studied, including several simultaneous multi-wavelength observation campaigns (see, e.g., \cite{Example_1553_MWL}). However, the only data in the high-energy hump are from HESS and MAGIC. The measured VHE spectra are unusually soft (photon index $\Gamma$=4.0$\pm$0.6 and $\Gamma$=4.2$\pm$0.3, respectively) but the errors are large, clearly requiring improved measurements before detailed interpretation of the complete SED is possible. Further complicating any SED interpretation is the absorption of VHE photons (\cite{EBL_effect3,EBL_effect2}) by pair-production on the Extragalactic Background Light (EBL). This absorption, which is energy dependent and increases strongly with redshift, distorts the VHE energy spectra observed from distant objects. For a given redshift and a given EBL model, the effects of the latter on the observed spectrum can be reasonably accounted for during SED modeling. Unfortunately, the redshift of PG\,1553+113 is currently unknown\footnote{\cite{redshift_question} demonstrate the catalog redshift of $z=0.36$ is incorrect.}. To date no emission or absorption lines have been measured from PG\,1553+113 despite more than ten observation campaigns with optical instruments, including EMMI at the NTT and FORS2 with the 8-meter VLT telescopes (\cite{Carangelo_03,no_lines}). Lower limits of $z>0.09$ (\cite{no_lines}) and $z>0.3$ (\cite{Carangelo_03}) were determined from the lack of detected absorption/emission lines, implying that the effect of the EBL is large in the observed VHE data. The absence of absorption and emission lines suggests that the non-thermal component of the emission from PG\,1553+113 is largely dominant over that of the host galaxy. This is consistent with the fact that, although some hints for a host galaxy have been suspected, no clear detection has yet been found, even in Hubble Space Telescope (HST) images of PG\,1553+113 taken during the HST survey of 110 BL Lac objects \cite{Hubble_image}. Interestingly, approximately 80\% (88/110) of the BL\,Lacs initially surveyed by HST have known redshifts, of which all 39 with $z<0.25$ and 21 of the 28 with $0.25<z<0.6$ have their hosts resolved, suggesting that the redshift of PG\,1553+113 is indeed large. Recently, using a re-analysis of the HST snapshot survey of BL\,Lacs, \cite{z_extreme} claimed the dispersion of the absolute magnitude of BL\,Lac hosts is sufficiently small that the measurement of the host-galaxy brightness allows a reliable estimate of their redshift. With the assumption that there is no strong evolution, these authors set a possible lower limit of $z>0.78$ for PG\,1553+113. However, adding new STIS-HST data of $z>0.6$ BL\,Lacs to the snapshot survey, \cite{HST_new} found on the contrary that host galaxies of BL\,Lacs evolve strongly, making any photometric redshift determination questionable. A conservative upper limit ($z<0.74$) was determined (\cite{HESS_discovery}) from the photon spectrum measured by HESS. The same limit was similarly determined from the MAGIC spectral measurement (\cite{MAGIC_1553}). Using both the MAGIC and HESS data, a stronger upper limit of $z<0.42$ was later reported \cite{Mazin_limit} based on the assumption that there is no break in the intrinsic VHE spectrum of the object. The range of allowed redshift is clearly large enough that the significant effects of EBL absorption cannot be reliably removed from the observed VHE spectrum of PG\,1553+113. This correspondingly makes modeling the high-energy portion of its SED unreliable. A clear detection of the object's redshift would dramatically improve the understanding of PG\,1553+113. Further, if it were found to be distant, its VHE spectrum could potentially provide strong constraints on the poorly-measured EBL (assuming a reasonable intrinsic spectrum) and contribute to establishing the VHE $\gamma$-ray horizon. Results from 17.2 hours of new HESS observations of PG\,1553+113 in 2006 are reported here. In addition, a re-analysis of the previously published 2005 HESS data (7.6 hours), with an improved calibration of the absolute energy scale of the detector, is presented. HESS and the Suzaku X-ray satellite observed the blazar simultaneously in July 2006. The HESS results from this epoch are also discussed. This will enable, for the first time, future modeling of an SED determined from simultaneous observations at VHE and lower energies. Finally, results of a March 2006 VLT SINFONI spectroscopy campaign to determine the redshift of PG\,1553+113 are also reported. | With a data set that is $\sim$3 times larger than previously published (\cite{HESS_discovery}), the HESS signal from PG\,1553+113 is now highly significant ($\sim$10$\sigma$). Thus, the evidence for VHE emission previously reported is clearly verified. However, the flux measured in 2005 is now $\sim$3 times higher than first reported due to an improved calibration of the absolute energy scale of HESS. The statistical error on the VHE photon index is now reduced from $\sim$0.6 to $\sim$0.3. Nevertheless, the error of 0.34 is still rather large, primarily due to the extreme softness of the observed spectrum ($\Gamma=4.46$). The total HESS exposure on PG\,1553+113 is $\sim$25 hours. Barring a flaring episode, not yet seen in two years of observations, a considerably larger total exposure ($\sim$100 hours) would be required to significantly improve the spectral measurement. This large exposure is unlikely to be quickly achieved. However, the VHE flux from other AGN is known to vary dramatically and even a factor of a few would reduce the observation requirement considerably. Should such a VHE flare occur, not only will the error on the measured VHE spectrum be smaller, but the measured photon index may also be harder (see, e.g., \cite{VHE_hardening}). Both effects would dramatically improve the redshift constraints and correspondingly the accuracy of the source modeling. Therefore, the VHE flux from PG\,1553$+$113 will continue to be monitored by HESS. In addition, the soft VHE spectrum makes it an ideal target for the lower-threshold HESS Phase-II \cite{HESSII} which should make its first observations in 2009. | 7 | 10 | 0710.5740 |
0710 | 0710.2160_arXiv.txt | We have searched for pulsed radio emission from magnetar 4U 0142+61 at the frequency of 111 MHz. No pulsed signal was detected from this source. Upper limits for mean flux density are 0.9~-~9~mJy depending on assumed duty cycle (.05~-~.5) of the pulsar. | 4U~0142+61 is an anomalous X-ray pulsar (AXP) with 8.7 seconds period. At period derivative $\dot P = 2 \times 10^{-12}$ magnetic field on a surface of a neutron star equals to $1.3 \times 10^{+14}$~G - typical value for magnetars. The pulsar can be seen in hard X-ray (\cite{den06}, \cite{kui06}), optical (\cite{ker02}), and infrared (\cite{hul04}, \cite {wan06}) bands. We searched for pulsed radio emission of this AXP at the frequency of 111~MHz. | The search of pulsed radio emission from anomalous X-ray pulsar (magnetar) 4U~0142+61 at the frequency of 111~MHz give no positive results. Upper limits for mean flux density are 0.9~-~9~mJy depending on assumed duty cycle (.05 - .5) of the pulsar. | 7 | 10 | 0710.2160 |
0710 | 0710.2210_arXiv.txt | We study the evolution of simple cosmic string loop solutions in an inflationary universe. We show, for the particular case of circular loops, that periodic solutions do exist in a de Sitter universe, below a critical loop radius $R_c H=1/2$. On the other hand, larger loops freeze in comoving coordinates, and we explicitly show that they can survive more $e$-foldings of inflation than point-like objects. We discuss the implications of these findings for the survival of realistic cosmic string loops during inflation, and for the general characteristics of post-inflationary cosmic string networks. We also consider the analogous solutions for domain walls, in which case the critical radius is $R_c H=2/3$. | Introduction} Topological defects are unavoidably formed at cosmological phase transitions \cite{Kibble,Book}. Studying their physical properties, evolution and cosmological consequences is therefore mandatory for a proper understanding of the early universe. The last three decades have seen dramatic progress in this task (see \cite{Book} for a review), but significant knowledge gaps still exist. The aim of the present paper is to eliminate one of these gaps. Defect-forming phase transitions often occur near or at the end of inflation. Moreover, it is possible that various stages of inflation occur, with defects being formed in between them \cite{openinf}. It is therefore important to understand the effects of inflation on the defects, as well as to quantify their ability to survive any inflationary periods that occur after they form. The effects of inflation on the internal (microscopic) structure of defects have been studied in \cite{basu}, which shows that those with thickness $\delta H>1/\sqrt{2}$ get smeared by expansion, while those with smaller thickness survive. Effects on macroscopic (cosmological) scales are well known for the case of long string networks \cite{quantitative,extending}, but this is not so for the loop populations, despite the fact that they are known to contain a total amount of energy which is comparable to (if not greater than) that in the long strings \cite{quantitative,protyloops}. Here we study the effects of inflation on specific loop solutions, and consider both their microscopic and macroscopic evolution. We also discuss the consequences of our findings for cosmological scenarios involving string networks. Incidentally, we note that loops in a flat anisotropic universe were studied in \cite{anisot}, where it was shown that the anisotropy of the background has an effect on the loop motion; the reasons for this will become clearer in what follows. We will also present a brief analysis of the analogous solutions for domain walls. We shall assume that the source of inflation is a perfect fluid with equation of state $p=w\rho$ (with $ w < -1/3$) and the scale factor behaves as $a \propto t^{2/3(1+w)}$; if $w = -1$ then $a \propto \exp(Ht)$ with $H$ being constant. | 7 | 10 | 0710.2210 |
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0710 | 0710.5430_arXiv.txt | \footnotesize\ This paper\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. 143-156} addresses the problem of scattering line polarization and the Hanle effect in one-dimensional (1D), two-dimensional (2D) and three-dimensional (3D) media for the case of a two-level model atom without lower-level polarization and assuming complete frequency redistribution. The theoretical framework chosen for its formulation is the QED theory of Landi Degl'Innocenti (1983), which specifies the excitation state of the atoms in terms of the irreducible tensor components of the atomic density matrix. The self-consistent values of these density-matrix elements is to be determined by solving jointly the kinetic and radiative transfer equations for the Stokes parameters. We show how to achieve this by generalizing to Non-LTE polarization transfer the Jacobi-based ALI method of Olson {\it et al.} (1986) and the iterative schemes based on Gauss-Seidel iteration of Trujillo Bueno and Fabiani Bendicho (1995). These methods essentially maintain the simplicity of the $\Lambda-$iteration method, but their convergence rate is extremely high. Finally, some 1D and 2D model calculations are presented that illustrate the effect of horizontal atmospheric inhomogeneities on magnetic and non-magnetic resonance line polarization signals. | The scattering line polarization, and its modification due to a weak magnetic field ---such that the Zeeman splitting is negligible compared with the line width (the so called Hanle effect; Hanle, 1924)---, sensitively depends on the {\it anisotropy} of the radiation field and on the magnetic field vector geometry (Landi Degl'Innocenti, 1985; Stenflo, 1994). The solar atmospheric plasma is spatially inhomogeneous, with vertical and horizontal variations, not only in temperature, macroscopic velocity and density, but also in the orientation and intensity of the magnetic field (see S\'anchez Almeida, 1999 for new insights in this respect). It is thus clear that, in order to fully exploit the Hanle effect as a diagnostic tool for weak magnetic fields, we also need to address the problem of resonance line polarization and the Hanle effect in 2D and 3D media where the radiation field's anisotropy is different from that corresponding to the currently-assumed 1D atmospheric models. To this end, this contribution begins presenting a formulation of resonance line polarization and the Hanle effect that we consider as the most suitable one for practical RT applications. It is based on the density-matrix theory for the generation and transfer of polarized radiation (see Landi Degl'Innocenti, 1983; 1984; 1985). In this paper we consider the standard case of a two-level model atom neglecting atomic polarization in its lower level (i.e. it is assumed that the lower-level Zeeman sublevels are equally populated and that there are no coherences among them). The quantities whose {\it self-consistent} values are to be determined are the irreducible tensor components of the density matrix ($\rho^K_Q$), which depend only on the spatial coordinates. The statistical equilibrium (SE) and RT equations to be solved are valid independently of whether we assume 1D, 2D or 3D geometries. A summary of previous work done in the subject of the numerical solution of Non-LTE polarization transfer problems can be found in Trujillo Bueno and Manso Sainz (1999). In this respect, we should mention the recent work of Nagendra {\it et al.} (1998; see also their contribution in these proceedings) where the Hanle effect in 1D is considered using a different theoretical approach. For information concerning the numerical solution of more general polarization transfer problems formulated with the density-matrix theory see Trujillo Bueno (1999). The outline of this paper is as follows. Section 2 presents the basic equations for the case of a two level atom without lower-level atomic polarization and assuming complete frequency redistribution. Section 3 shows how the very efficient iterative methods of solution investigated by Trujillo Bueno and Fabiani Bendicho (1995) (Jacobi, Gauss-Seidel and successive over-relaxation) can be suitably generalized to the problem of Non-LTE polarized radiative transfer in the Hanle effect regime. Finally, in Sect. 4 we present the results of some illustrative 2D Hanle-effect calculations for triplet lines and discuss the ensuing horizontal transfer effects.% | We have developed a Hanle effect code that allows the numerical simulation of resonance line polarization signals in the presence of weak magnetic fields in 1D, 2D and 3D media. The governing equations have been formulated working within the framework of the density matrix polarization transfer theory of Landi Degl'Innocenti (1983, 1985). These SE and RT equations are the same independently of whether we are considering 1D, 2D or 3D atmospheric models. The six $\rho^K_Q$-unknowns of the problem are neither frequency nor angle dependent, since they only vary with the spatial position. Three different iterative schemes that were originally developed for RT applications in the unpolarized case have been generalized to solve this set of equations: Jacobi (ALI), Gauss-Seidel and SOR (Trujillo Bueno \& Fabiani Bendicho, 1995, Trujillo Bueno \& Manso Sainz 1999). This kind of iterative methods does not make use of any matrix inversion, and essentially maintain the $\Lambda$-iteration simplicity. The only difference between the 1D, 2D and 3D versions of our Hanle effect code lies in the formal solution routine that calculates the radiation field tensors from the current values of the density-matrix elements. To this end, in 2D we use the formal solver developed by Auer, Fabiani Bendicho and Trujillo Bueno (1994) and in 3D we use the one presented at this workshop by Fabiani Bendicho and Trujillo Bueno (1999). We have also shown some Hanle effect results for 1D and 2D media. These calculations illustrate how weak magnetic fields and horizontal radiative transfer effects compete to modify the scattering line polarization signals expected from plane-parallel 1D atmospheres. Thus, further careful investigations must be done in order to separate both effects, when diagnosing weak solar magnetic fields via the Hanle effect. This type of future studies should be done thinking in the interpretation of {\it low} spatial resolution scattering line polarization observations. Our numerical approach is very efficient and suitable to investigate scattering polarization signals for a variety of atmospheric models having any desired temperature, density and magnetic field vector variations. Another useful research that can be done with our Hanle effect codes concerns the simulation of polarization signals emerging from realistic MHD and semi-empirical 2D and 3D models. | 7 | 10 | 0710.5430 |
0710 | 0710.5882.txt | {}{}{}{}{} % 5 {} token are mandatory \abstract % context heading (optional) % {} leave it empty if necessary {The young $\sigma$~Orionis cluster is an indispensable basis for understanding the formation and evolution of stars, brown dwarfs and planetary-mass objects. Our knowledge of its stellar population is, however, incomplete.} % aims heading (mandatory) {I present the Mayrit catalogue, that comprises most of the stars and high-mass brown dwarfs of the cluster.} % methods heading (mandatory) {The basis of this work is an optical-near infrared correlation between the 2MASS and DENIS catalogues in a circular area of radius 30\,arcmin centred on the OB-type binary $\sigma$~Ori~AB. The analysis is supported on a bibliographic search of confirmed cluster members with features of youth and on additional X-ray, mid-infrared and astrometric data.} % results heading (mandatory) {I list 241 $\sigma$~Orionis stars and brown dwarfs with known features of youth, {97} candidate cluster members (40 are new) and {115} back- and foreground sources in the survey area. The {338} cluster members and member candidates constitute the Mayrit catalogue.} % conclusions heading (optional), leave it empty if necessary {This catalogue is a suitable input for studying the spatial ditribution, multiplicity, properties and frequency of discs and the complete mass function of $\sigma$~Orionis.} | The $\sigma$ Orionis cluster in the \object{Ori~OB~1b} Association is getting as important for the study of the formation, evolution and characterisation of stars and substellar objects as other famous clusters and star-forming regions, like the \object{Hyades}, the \object{Pleiades}, the \object{Orion Nebula Cluster} or the \object{Taurus-Auriga Complex}. The $\sigma$~Orionis cluster is young (3$\pm$2\,Ma), nearby ( $\sim$385\,pc) and relatively free of extinction (Lee 1968; Brown et~al. 1994; Oliveira et~al. 2002; Zapatero Osorio et~al. 2002a; Sherry et~al. 2004; B\'ejar et~al. 2004b; Caballero 2007d). Firstly identified by Garrison (1967) and Lyng\aa~(1981), $\sigma$~Orionis was rediscovered by Wolk (1996) and Walter et~al. (1997). They reported a clustering of young low-mass stars, many of them positionally coincident with X-ray sources, surrounding the Trapezium-like, multiple stellar system $\sigma$~Ori in the vicinity of the \object{Horsehead Nebula} (see Caballero 2007b for a description of the multiple system $\sigma$~Ori and its surroundings). Previously, the area had been investigated during wide searches in the Orion complex with prism-objective and Schmidt plates, detecting a wealth of emission stars (e.g. Haro \& Moreno 1953; Wiramihardja et~al. 1989), but the cluster had not been treated as an independent entity within the complex. Complete compilations of the determinations in the literature of the age, heliocentric distance, frequency of discs and mass function of the $\sigma$~Orionis cluster are in Caballero (2007a). A chapter of the Handbook of Star Forming Regions, edited by B.~Reipurth, will be exclusively devoted to $\sigma$~Orionis (F.~M.~Walter et~al., in~press). After the seminal work by B\'ejar et~al. (1999), who found for the first time a rich population of young brown dwarfs in $\sigma$~Orionis, the cluster has turned into an excellent laboratory for the study of, e.g.: \begin{itemize} \item the search for free-floating planetary-mass objects (with masses below the deuterium-burning limit) and the study of the substellar mass function down to a few Jupiter masses (Zapatero Osorio et~al. 2000; B\'ejar et~al. 2001; Gonz\'alez-Garc\'{\i}a et~al. 2006; Caballero et~al. 2007); \item the frequency and the properties of $\sim$3\,Ma-old discs at different mass intervals (Jayawardhana et~al. 2003; Oliveira et~al. 2004, 2006; Hern\'andez et~al. 2007; Caballero et~al. 2007; Zapatero Osorio et~al. 2007a); \item the masses of OB-type stars in resolved binary systems (Heintz 1997; Mason et~al. 1998; Caballero 2007d); \item the X-ray emission of young stars and brown dwarfs (Mokler \& Stelzer 2002; Sanz-Forcada et~al. 2004; Franciosini et~al. 2006; Caballero 2007b); \item the characterisitics of jets and Herbig-Haro objects (Reipurth et~al. 1998; L\'opez-Mart\'{\i}n et~al. 2001; Andrews et~al. 2004); and \item the photometric variability of low-mass stars and brown dwarfs (Bailer-Jones \& Mundt 2001; Caballero et~al. 2004; Scholz \& Eisl\"offel~2004). \end{itemize} Many interesting star-like objects have been discovered in the cluster, from the helium-rich, B2.0Vp-type magnetic star $\sigma$~Ori~E (Greenstein \& Wallerstein 1958), through the Class~I object candidate IRAS~05358--0238 (Oliveira \& van Loon 2004), to the hypothetical proplyd $\sigma$~Ori~IRS1 (van Loon \& Oliveira 2003; Caballero 2005, 2007b). The most interesting objects in the cluster are, however, below the hydrogen-burning mass limit. Some of these substellar objects are the $\sim$T6-type object \object{S\,Ori~70} (which may be the least massive body directly detected out of the Solar System, $\sim$3\,$M_{\rm Jup}$ -- Zapatero Osorio et~al. 2002c, 2007b; Burgasser et~al. 2004), the T Tauri substellar analog S\,Ori~J053825.4--024241 (which is the most variable brown dwarf yet found; Caballero et~al. 2006a), the two strong H$\alpha$ emitters at the planetary boundary \object{S\,Ori~55} and \object{S\,Ori~71} (with masses of only 10--20\,$M_{\rm Jup}$ and equivalent widths of the H$\alpha$ line of up to --700\,\AA; Zapatero Osorio et~al. 2002b; Barrado y Navascu\'es et~al. 2002a) and the brown dwarf-exoplanet system candidate \object{SE~70} + \object{S\,Ori~68} (which could be the widest planetary system known so far; Caballero et~al. 2006b). The number of known substellar objects in $\sigma$~Orionis is comparable to those of other rich, more massive, younger star-forming regions like Chamaeleon, Ophiuchus or the Orion Nebula Cluster. However, $\sigma$~Orionis is by far the region with the largest amount of brown dwarfs with membership confirmation ($>$ 30; Caballero et~al. 2007) and planetary-mass object candidates (29; Zapatero Osorio et~al. 2000; Gonz\'alez-Garc\'{\i}a et~al. 2006; Caballero 2007b; Caballero et~al. 2007). Many of the latter bodies have measured L spectral types and/or flux excess longwards of 5\,$\mu$m (Zapatero Osorio et~al.~(2007a). Important efforts have been recently carried out to characterise the $\sigma$~Orionis cluster in general and to investigate the connection between its stellar and substellar populations in particular (e.g. B\'ejar et~al. 2004a; Sherry et~al. 2004; Kenyon et~al. 2005; Burningham et~al. 2005; Caballero 2005, 2007a, 2007b, 2007c). The works by Kenyon et~al. (2005), who investigated membership, binarity and accretion among very low-mass stars and brown dwarfs surrounding $\sigma$~Ori, and, especially, Sherry et~al. (2004) stand out. The latter authors presented an ambitious study, estimating the number of cluster members in the mass range 0.2\,$M_\odot$ $\lesssim M \lesssim$ 1.0\,$M_\odot$ and the radius, age and total mass of the cluster. All these works are, however, incomplete or biased in some way: no comprehensive, homogenous, multi-band study, fully covering the whole stellar mass interval (from $\sim$20\,$M_\odot$ down to the substellar boundary) and the cluster area without gaps (from the very centre to the border) exists so far in $\sigma$~Orionis. I extend the study of the brightest stars of the cluster shown in Caballero (2007a) down to the hydrogen-burning mass limit and beyond. I start out from a correlation between the 2MASS and DENIS catalogues within the environment of the Virual Observatory\footnote{See {http://www.ivoa.net}.} and a bibliographic search of confirmed cluster members, and support them with spectroscopic and photometric data from the X-ray region to 120\,$\mu$m when available. The outcome of this work is the Mayrit catalogue, that comprises the majority of the stars and high-mass brown dwarfs of $\sigma$~Orionis. | The $\sim$3\,Ma-old $\sigma$~Orionis cluster is a new cornerstone for observational and theoretical studies with the aim to understand the general processes of collapse and fragmentation of a molecular cloud, formation of stars and substellar objects and evolution of circumstellar discs. I~present a comprehensive, rather complete catalogue of cluster members that can be used for further studies in $\sigma$~Orionis (e.g. spatial distribution, multiplicity, mass function and frequency and characterisation of~discs). This investigation covers the whole stellar and part of the brown dwarf domain of the cluster from $\sim$18 to $\sim$0.03\,$M_\odot$. I~have performed an $IK_{\rm s}$ survey in a 30\,arcmin-radius region centred on the O9.5V+B0.5V binary $\sigma$~Ori~AB using the Aladin tool and optical and near-infrared data from the DENIS and 2MASS catalogues. The photometric data have been complemented with information from the literature regarding the membership of known sources. I~have compiled a list of 241 cluster stars and brown dwarfs with known features of youth (e.g. early spectral types, Li~{\sc i} in absorption, near and mid infrared excess attributed to surrounding discs, strong X-ray and H$\alpha$ emissions), 85 fore- and background stars with astrometric and spectroscopic data (nine are new), {18} galaxies with extended FWHMs and 12 probable reddened sources in a nebulosity northeastern of the survey area, associated to the Horsehead Nebula. From the $I$ vs. $I-K_{\rm s}$ diagram, I have selected {97} additional photometric cluster member candidates without reliable membership information. This makes a list of {338} $\sigma$~Orionis members and member candidates, the Mayrit catalogue, of which more than 70\,\% display features of extreme youth. %Global contamination and incompleteness are at the order of magnitude of %only 10\,\%. I~tabulate precise coordinates, $IJHK_{\rm s}$ and suplementary information of all the cluster members, member candidates and non-members. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% | 7 | 10 | 0710.5882 |
0710 | 0710.3558.txt | We investigate the dynamics of magnetic fields in spiral galaxies by performing 3D Magnetohydrodynamics (MHD) simulations of galactic discs subject to a spiral potential using cold gas, warm gas and a two phase mixture of both. Recent hydrodynamic simulations have demonstrated the formation of inter-arm spurs as well as spiral arm molecular clouds provided the ISM model includes a cold HI phase. We find that the main effect of adding a magnetic field to these calculations is to inhibit the formation of structure in the disc. However, provided a cold phase is included, spurs and spiral arm clumps are still present if $\beta \gtrsim 0.1$ in the cold gas. A caveat to the two phase calculations though is that by assuming a uniform initial distribution, $\beta \gtrsim 10$ in the warm gas, emphasizing that models with more consistent initial conditions and thermodynamics are required. Our simulations with only warm gas do not show such structure, irrespective of the magnetic field strength. Furthermore, we find that the introduction of a cold HI phase naturally produces the observed degree of disorder in the magnetic field, which is again absent from simulations using only warm gas. Whilst the global magnetic field follows the large scale gas flow, the magnetic field also contains a substantial random component that is produced by the velocity dispersion induced in the cold gas during the passage through a spiral shock. Without any cold gas, the magnetic field in the warm phase remains relatively well ordered apart from becoming compressed in the spiral shocks. Our results provide a natural explanation for the observed high proportions of disordered magnetic field in spiral galaxies and we thus predict that the relative strengths of the random and ordered components of the magnetic field observed in spiral galaxies will depend on the dynamics of spiral shocks. | Observations of the magnetic field in galaxies indicate that the magnetic pressure is comparable in magnitude to the thermal pressure and interstellar turbulence \citep{Heiles2005}. Consequently magnetic fields are expected to play an important part in the dynamics of the ISM, including spiral shocks and the formation of molecular clouds. The response of a gaseous disc to a spiral potential has been well examined both theoretically and numerically \citep*{Roberts1969,Shu1972,Woodward1976,Gittins2004,Wada2004,DBP2006}. The gas flow evolves into a quasi-stationary solution which contains spiral shocks, providing the potential is of sufficient strength. \citet{Roberts1970} found that the strength of the shock decreases with increasing magnetic field strength, whilst the magnetic field strength is amplified in the spiral shock, the latter agreeing with observations at the time for galactic magnetic fields. There are numerous hydrodynamic simulations of gas discs subject to a spiral potential (e.g. \citealt{Kim2002,Chak2003,Slyz2003,Wada2004,Gittins2004,Dobbs2006}) which discuss the formation of substructure and location of the shock. Whilst 2D simulations of warm gas show the formation of small-scale spurs perpendicular to the spiral arms \citep{Gittins2004,Wada2004}, these features appear to be suppressed in 3D simulations \citep{Kim2006}. On the other hand, \citet{Dobbs2006} do find more extensive substructure and spiral arm clouds \citep*{DBP2006,Dobbs2007} in 3D simulations, but only in models where the gas temperature is $\leq 1000$ K. There are comparatively fewer MHD calculations, largely due to the relatively recent addition of magnetic fields into numerical codes, and the limitations of numerical resolution. \citet{Shetty2006} investigate the formation of spurs using 2D MHD grid-based calculations of the warm ISM with self gravity, concluding that self gravity is necessary for spurs to evolve. \citet{Gomez2002,Gomez2004} also perform 3D calculations with warm gas, but whilst both groups compare MHD with non-MHD results, neither investigate the effects of varying the field strength or ISM temperature. Recent observations now provide much more detailed measurements of the magnetic field strength and direction in the diffuse (warm) ISM (see reviews by \citealt{Beck1996} and \citealt{Beck2007}). Spiral magnetic arms appear to occur in all disc galaxies, regardless of the presence of optical spiral arms \citep{Beck2005}. Typically the field in the inter-arm regions is several $\mu$G, but may be 10 or more $\mu$G in the spiral arms, and the total field contains random and ordered (regular) components of comparable strength. The general consensus currently favours dynamo theory to explain the origin of these relatively strong magnetic fields \citep*{Parker1971,P1971,Balsara2004,Beck2005}, whereby the magnetic field is generated by turbulence in the ISM, and can produce or enhance magnetic spiral arm arms \citep*{Panesar1992,Rohde1998}. However in grand design galaxies where the gas experiences strong spiral shocks, such as M51, the magnetic spiral arms tend to be aligned with the dust lanes \citep{Nein1992} suggesting that the magnetic field is strongly related to the dynamics of the spiral shock. Magnetic fields have been long associated with molecular cloud formation, originally through the Parker instability \citep{Parker1966}, and more recently via the magneto-rotational \citep*{Kim2003,Piontek2005} instability. The Parker instability is expected to produce sinusoidal motions in the $z$ direction, with the magnetic field channeling gas into dense concentrations in the plane of the disc. Molecular clouds formed in this way were originally believed to be magnetically supported, with lifetimes of order $10^8$ years \citep*{Zweibel1975,Shu1987}. However the Parker instability is found to be relatively weak, and singularly insufficient to induce molecular cloud formation \citep*{Elmegreen1982,OtherKim2001,KOS2002}. The magneto-rotational instability (MRI) has been proposed as a formation mechanism away from spiral arms \citep{Kim2003}. However the MRI generally takes a few orbits to initialise, particularly if the field is toroidal (e.g. 10's of orbits in simulations by \citealt{Nish2006}), and is therefore of less importance where spiral shocks occur. Magnetic fields may be more relevant to GMC (giant molecular cloud) formation in conjunction with gravitational instabilities \citep{ElmegreenMJI1987,Kim2001,KOS2002} [also known as the magneto-Jeans-instability, MJI] or cooling \citep{Kosinski2007}. In this paper we describe 3D MHD calculations of a galactic disc subject to a spiral perturbation. These are the first fully Lagrangian MHD calculations to model this problem, using the Smoothed Particle Magneto-hydrodynamics (SPMHD) code \citep{Price2004,Price2005,PB2007}. We compare the structure of the disc for a range of magnetic field strengths, assuming a single phase medium of cold or warm gas, or a two-phase medium. We also describe the strength and morphology of the magnetic field, and relate these to observations. In this paper we focus on the effect of the spiral shock inducing structure in the gas rather than Parker or MRI instabilities, and we leave a discussion of results including self-gravity to a future paper. | We have performed simulations of galactic discs subject to a spiral potential with a range of field strengths. The main results we have discussed are 1) the reduction in structure across the disc as the magnetic field strength increases, and 2) the possibility of spiral shocks inducing an irregular magnetic field in the ISM. As the strength of the magnetic field increases, the strength of the spiral shocks and therefore density of the spiral arms are reduced. This is as expected from the analysis of \citet{Roberts1970}. Consequently the formation of spurs is increasingly suppressed for higher magnetic field strength. The spiral arms themselves are more continuous and clumps along the spiral arms are less dense. This supports previous 2D results \citep{Shetty2006,Tanaka2005} which find that both Kelvin-Helmholtz and gravitational instabilities perpendicular to the magnetic field are reduced by stronger magnetic fields (although we do not relate the structure in our simulations to these instabilities, our overall conclusions agree). We find that the addition of the magnetic field is similar in effect to an increase in thermal pressure, in that both provide a pressure which oppose the formation of structure, and smooth out the gas. Nonetheless, we still find significant inter-arm structure with the presence of a magnetic field, unlike the results (those which are non-self gravitating) of \citet{Shetty2006}. This structure is present in the cold component of the ISM, which has generally not been included in previous simulations. Inhomogeneities present in the initial random distribution of gas become amplified by spiral shocks. With warm gas or strong magnetic fields, these inhomogeneities are smoothed out by the pressure. For our calculations with cold gas, we find that magnetic fields only prevent substructure when the ratio of gas to magnetic pressure ($\beta$) is $\lesssim 0.1$. For $\beta \gtrsim 1$, spurs perpendicular to the arm still form in the cold gas. The additional pressure provided by warm gas in the two-phase results increases the longevity of structure in the inter-arm regions, even when the magnetic field dominates for the cold gas. The main difference between ours and previous results is that we use cold gas, although we also use a weaker magnetic field than \citet{Shetty2006}. We find structure is present for values of $\beta$ in the cold gas similar to observations, but the volume averaged magnetic field strengths in our two phase models are typically $< 1 \mu$G, lower than observations. This discrepancy arises because we start with an initially uniform distribution of cold gas. To obtain a similar $\beta$ for both the cold and warm gas in our two phase models, we would need the cold gas to be 100 times denser than the warm gas, i.e. this just indicates that structure is \emph{already} present in the cold gas. If we used higher field strengths, we would find, like the results of \citet{Shetty2006}, very little structure emerges. In this case, $\beta$ for the warm gas would be closer to observations, but the magnetic pressure would dominate the thermal pressure by a factor of 100 or so in the cold gas, inconsistent with observations. Thus we either need to distribute the cold gas initially in clumps, or ideally include a much more detailed thermal treatment of the ISM, in order to obtain magnetic field strengths \emph{and} $\beta$ consistent with observations. We find that whilst the ISM appears highly structured in observations, we would expect a higher degree of structure with relatively weak magnetic fields. For instance, the two phase model where $\beta_{cold}=4$ retains much more structure typical of grand design galaxies compared to the case where $\beta_{cold}=0.4$. Current observations suggest that magnetic pressure exceeds thermal pressure \citep{HT2005}. We however note that $\beta$ exhibits a range of values in our simulations, and $\beta$ tends to be take smaller values in the spiral arms and dense gas (for a given temperature), which are more likely to correspond with observations. Again a more complete treatment of the ISM in future work, and further observations of the CNM will allow a better comparison. In our simulations, the magnetic field is compressed by the spiral shocks, as expected from a straightforward analysis of MHD shocks \citep{Roberts1970,Priest1982}. The relative increase in the magnetic field strength is greater where the shock is stronger. However the most intriguing result from our simulations is the possibility that spiral shocks generate an irregular magnetic field. This process has not been identified in previous simulations, which we attribute to the fact that they have not included a cold phase. Galaxies are known to contain a random component of the field, but it is usually supposed that this is due to supernovae and/or feedback from stars. We therefore postulate that spiral shocks are important in generating disorder in the magnetic field, whilst simultaneously inducing a velocity dispersion and density structure in the gas (\citealt*{Bonnell2006,KKO2006}; \citealt{DBP2006}). The degree of order in the disc, similar to the presence of arm/inter-arm structure is related to the strength of the shock. When the gas is cold and the magnetic field is weak, the shock is much stronger, and a higher velocity dispersion is induced in the gas. Consequently the magnetic field lines follow the gas and the field becomes tangled. In simulations with warm gas, we find the magnetic field is almost entirely regular, although with a stronger shock, the field may become more irregular. However in our two-phase results, the velocities and irregularities in the magnetic field of the cold gas induces comparatively more disorder in the warm gas than the single phase calculations. Although there are no measurements of regular versus disordered components of the magnetic field in cold HI, we would predict from our results that the field in the cold gas will be more disordered than that of warm HI. | 7 | 10 | 0710.3558 |
0710 | 0710.1554_arXiv.txt | {}{We have investigated a sample of 28 well-known spectroscopically-identified magnetic Ap/Bp stars, with weak, poorly-determined or previously undetected magnetic fields. The aim of this study is to explore the weak part of the magnetic field distribution of Ap/Bp stars.} {Using the MuSiCoS and NARVAL spectropolarimeters at T\'elescope Bernard Lyot (Observatoire du Pic du Midi, France) and the cross-correlation technique Least Squares Deconvolution (LSD), we have obtained 282 LSD Stokes $V$ signatures of our 28 sample stars, in order to detect the magnetic field and to infer its longitudinal component with high precision (median $\sigma=40$~G). } {For the 28 studied stars, we have obtained 27 detections of Stokes $V$ Zeeman signatures from the MuSiCoS observations. Detection of the Stokes $V$ signature of the $28^{\rm th}$ star (HD~32650) was obtained during science demonstration time of the new NARVAL spectropolarimeter at Pic du Midi. This result shows clearly that when observed with sufficient precision, all firmly classified Ap/Bp stars show detectable surface magnetic fields. Furthermore, all detected magnetic fields correspond to longitudinal fields which are significantly greater than some tens of G. To better characterise the surface magnetic field intensities and geometries of the sample, we have phased the longitudinal field measurements of each star using new and previously-published rotational periods, and modeled them to infer the dipolar field intensity and the magnetic obliquity. The distribution of derived dipole strengths for these stars exhibits a plateau at about 1 kG, falling off to larger and smaller field strengths. Remarkably, in this sample of stars selected for their presumably weak magnetic fields, we find only 2 stars for which the derived dipole strength is weaker than 300~G. We interpret this ``magnetic threshold'' as a critical value necessary for the stability of large-scale magnetic fields, and develop a simple quantitative model that is able to approximately reproduce the observed threshold characteristics. This scenario leads to a natural explanation of the small fraction of intermediate-mass magnetic stars. It may also explain the near-absence of magnetic fields in more massive B and O-type stars.}{}{} | The magnetic chemically peculiar Ap/Bp stars are the non-degenerate stars for which the strongest magnetic fields have been measured (Landstreet 1992). Although the fields are thought to be fossil remnants of flux swept up during star formation or produced via dynamo action on the pre-main sequence, their origin is not understood in any real detail (e.g. Moss 2001). Furthermore, the role of the magnetic field in the diffusion processes which are responsible for their chemical peculiarity has been studied in only a schematic fashion. Although more than one thousand main sequence {A-type stars} have been catalogued as magnetic Ap/Bp stars (Renson et al. 1991) following the scheme of Preston (1974), direct measurements of the magnetic field have been obtained for only a few hundred of them (Romanyuk 2000, Bychkov et al. 2003). Examination of the published measurements shows that the majority of the reported values are rather large. For example, $55\%$ of the 210 stars of the catalogue of Romanyuk (2000) with published magnetic field measurements have a maximum unsigned line-of-sight (longitudinal) magnetic field $B_\ell$ larger than 1 kG. On the other hand, according to Bohlender \& Landstreet (1990), the median root-mean-square (rms) longitudinal magnetic field of Ap stars (based on a small magnitude-limited sample observed by Borra \& Landstreet 1980) is only about 300~G (the largest rms field they report is only 710 G). This implies that most Ap stars have relatively weak ($\ltsim 1$~kG) magnetic fields, and that the available observations are strongly biased toward stars with the strongest and most easily-measured fields. One consequence of this bias is that the weak-field part of the magnetic field distribution of Ap stars is poorly studied. It is not known if it increases monotonically toward arbitrarily small field strength, or if it is truncated at a minimum magnetic field strength (as proposed by Glagolevskij \& Chountonov 2002). In order to improve our knowledge of the weak-field part of the magnetic field distribution of Ap stars, we have undertaken a study of a sample of 28 well-known spectroscopically-identified Ap/Bp stars, with very weak, poorly-determined or previously-undetected magnetic fields. We describe our survey in Sect. 2 and report our observational and modeling results in Sect. 3 and Sect. 4. We discuss the implications of our results, suggesting one possible interpretation involving the stability of large scale magnetic fields, in Sect. 5 and give our conclusions in Sect. 6. \section {The weak-field Ap stars survey} \subsection{The selected sample} Our sample is composed of spectroscopically-identified Ap/Bp stars belonging to the HD catalogue. Twelve stars were selected based on the observations of Borra \& Landstreet (1980) and Bohlender et al. (1993), identifying stars for which no significant detection of the magnetic field was obtained. Thirteen additional targets are stars for which only old photographic measurements of the magnetic field were available, typically by Babcock (1958), and for which measurements have poor precision and do not provide a significant detection of the magnetic field; these stars were generally selected using the catalogues of Romanyuk (2000) or Bychkov et al. (2003). Finally, 3 stars of our sample were selected which had not been observed for magnetic field before this work. The observational properties of the 28 stars are presented in Table 1. Section 3.2 gives more details on each star and on the obtained results. Because all of these stars are relatively bright, most have been known for decades and have been well studied. All appear in the Hipparcos catalogue (Perryman et al., 1997), and all but 6 have $\sigma_\pi/\pi<0.2$. The majority have photometrically-determined rotational periods and published values of $v\sin i$. Many have been studied using high-resolution spectroscopy and Doppler Imaging. Therefore the classification of this sample as {\em bona fide} Ap/Bp stars is generally quite firm. As a consequence of this careful selection (and as will be described later in the paper), {\em no} non-Ap/Bp star has been mistakenly included in the sample. \subsection{Observations and reduction} Stokes $V$ and Stokes $I$ spectra of the 28 sample stars were obtained during 10 observing runs, from July 2001 to June 2006. We used the MuSiCoS spectropolarimeter attached the Bernard Lyot telescope (TBL) at Observatoire du Pic du Midi. The MuSiCoS spectropolarimeter is composed of a cross-dispersed echelle spectrograph (Baudrand \& B\" ohm 1992) and a dedicated polarimeter module (Donati et al. 1999). The spectrograph is a table-top instrument, fed by a double optical fibre directly from the Cassegrain-mounted polarimeter. In one single exposure, this apparatus allows the acquisition of a stellar spectrum in a given polarisation (Stokes $V$ in this case) throughout the spectral range 450 to 660 nm with a resolving power of about 35000. Spectra in both orthogonal polarisations are recorded simultaneously by the CCD detector. A complete Stokes $V$ exposure consists of a sequence of four subexposures, between which the quarter-wave plate is rotated by 90$^{\rm o}$. This has the effect of exchanging the beams in the whole instrument, and in particular switching the positions of the two orthogonally polarised spectra on the CCD, thereby reducing spurious polarisation signatures. The echelle polarisation spectra were reduced using the ESpRIT package (Donati et al. 1997). The observation and reduction procedures are more thoroughly described by Shorlin et al. (2002). The correct operation of the MuSiCoS instrument, and in particular the absence of spurious magnetic field detections, is supported by other data obtained during these same observing runs, including studies of non-magnetic A-type stars (e.g. Shorlin et al. 2002), magnetic A, B and O-type stars (e.g., Ryabchikova et al. 2005a, Donati et al. 2001, Wade et al. 2006a) and magnetic late-type stars (e.g. Petit et al. 2005). MuSiCoS has recently been decomissioned, and has been replaced with NARVAL (Auri\`ere 2003), the new-generation spectropolarimeter which is a copy of the ESPaDOnS instrument in operation at the Canada-France-Hawaii Telescope (Donati 2004, 2007: in preparation). The main improvements of NARVAL in polarisation mode with respect to MuSiCoS are a spectral resolution of about 65000, spectral response between 370 nm and 1000 nm, and an overall sensitivity increased by a factor of about 30. \begin{table*}[t] \caption{Observational properties of the weak-field Ap star sample. Columns give ID and HD number, visual magnitude, spectral classification, effective temperature, luminosity and radius (with associated $1\sigma$ error bars), adopted LSD mask temperature, number of observations obtained and detection level (d=definite detection; m=marginal detection), maximum observed unsigned longitudinal field in G, $1\sigma$ error in G, and peak longitudinal field detection significance {\it z } = $ |B^{max}_{l}|$ / $\sigma$. $B_{\rm d}^{\rm min, 3.3}$ is the minimum dipole field (at $2\sigma$) inferred from the maximum measured longitudinal field and Eq. (7).} \begin{tabular}{lcrr|ccc|cccccc} \\ \hline ID & HD & $m_V$ & Spec & $T_{\rm eff}$ & $\log L$ & $R$ & Mask & \# & Det.& $|B_{\ell}|^{max}\pm \sigma$ & {\it z } & $B_{\rm d}^{\rm min, 3.3}$\\ & & & Type & (K) & ($L_\odot$) & ($R_\odot$) & (kK) & & level & (G)& &(G)\\ \hline \\ HN And& 8441 & 6.7 &A2p &$ 9060\pm 300$ & $ 1.90\pm 0.16$&$3.6 \pm 0.9 $ & 9 & 8& 8d & $157\pm 18 $& 8.7 & 399 \\ 43 Cas& 10221 & 5.5 &A0sp &$10660\pm 350$ & $ 2.11\pm 0.09$&$3.3 \pm 0.7 $ & 11 & 10& 8d & $148\pm 34 $& 4.3 & 264 \\ $\iota$ Cas &15089 & 4.5 &A5p &$ 8360\pm 275$ & $ 1.38\pm 0.05$&$2.3 \pm 0.4 $ & 9 & 12& 12d & $486\pm 23 $& 20.8& 1452 \\ &15144 & 5.8 &A6Vsp &$ 8480\pm 280$ & $ 1.21\pm 0.07$&$1.9 \pm 0.3 $ & 9 & 6& 6d & $631\pm 15 $& 42.0& 1983 \\ 21 Per& 18296 & 5.0 &B9p &$ 9360\pm 310$ & $ 2.08\pm 0.11$&$4.2 \pm 1.0 $ & 10 & 2& 2d & $213\pm 20 $& 10.6& 571 \\ 9 Tau & 22374 & 6.7 &A2p &$8390\pm 275$ & $ 1.48\pm 0.13$&$2.6 \pm 0.7 $ & 9 & 2& 2d & $523\pm 24 $& 21.7& 1568 \\ 56 Tau &27309 & 5.3 &A0sp &$12730\pm 420$ & $ 2.06\pm 0.08$&$2.2 \pm 0.4 $ & 12 & 12& 12d & $804\pm 50 $& 16.1& 2323 \\ 11 Ori& 32549 & 4.7 &A0sp &$10220\pm 335$ & $ 2.35\pm 0.12$&$4.7 \pm 1.2 $ & 11 & 11 & 1d3m & $186 \pm 39$& 4.7 & 356 \\ &32650 & 5.4 &B9sp &$11920\pm 390$ & $ 2.11\pm 0.07$&$2.7 \pm 0.5 $ & 12 & 18 & 2m3d& $91 \pm 18$& 5.0 & 237 \\ &37687 & 7.0 &B8 &$9450\pm 310$ & $ 2.18\pm 0.25$&$4.6 \pm 1.9 $ & 10 & 2 & 2d & $766\pm 119$& 6.4 & 1742 \\ 137 Tau& 39317 & 5.5 &B9spe &$10130\pm 330$ & $ 2.19\pm 0.14$&$4.0 \pm 1.1 $ & 11 & 8 & 3d1m & $216 \pm 59$& 3.6 & 323 \\ &40711 & 8.5 &A0 &$8070\pm 265$ & $ 1.94\pm 0.61$&$4.8 \pm 5.3 $ & 9 & 3 & 1d1m & $528 \pm 38$& 13.8& 1492 \\ &43819 & 6.2 &B9IIIsp &$10880\pm 355$ & $ 2.15\pm 0.20$&$3.3 \pm 1.2 $ & 11 & 8 & 8d & $628\pm 25 $& 25.1& 1907 \\ 15 Cnc& 68351 & 5.6 &B9sp &$10290\pm 340$ & $ 2.65\pm 0.21$&$6.6 \pm 2.4 $ & 10 & 16& 1d4m & $325 \pm 47$& 6.9 & 762 \\ 3 Hya &72968 & 5.7 &A1spe &$9840\pm 320$ & $ 1.55\pm 0.08$&$2.0 \pm 0.4 $ & 10 & 13& 13d & $427\pm 16 $& 27.4& 1304 \\ 45 Leo& 90569 & 6.0 &A0sp &$10250\pm 335$ & $ 1.78\pm 0.10$&$2.5 \pm 0.5 $ & 11 & 10& 10d & $541\pm 23 $& 23.5& 1634 \\ &94427 & 7.3 &A5 &$7250\pm 240$ & $ 1.05\pm 0.11$&$2.1 \pm 0.5 $ & 8 & 8& 8d & $356\pm 41 $& 8.6 & 904 \\ EP Uma &96707 & 6.0 &F0sp &$7780\pm 255$ & $ 1.54\pm 0.08$&$3.2 \pm 0.6 $ & 8 & 21& 7d7m & $69 \pm 33$& 2.3 & 128 \\ 65 Uma& 103498 & 6.9 &A1spe &$9220\pm 300$ & $ 2.06\pm 0.20$&$4.2 \pm 1.5 $ & 9 & 14& 12d1m & $169 \pm 19$& 8.9 & 432 \\ 21 Com& 108945 & 5.4&A2pvar&$8870\pm 290$ & $ 1.72\pm 0.09$&$3.1 \pm 0.6 $ & 9 & 13& 12d1m & $234 \pm 54$& 4.3 & 416 \\ $\omega$ Her & 148112&4.6&B9p&$9330\pm 305$ & $ 1.86\pm 0.08$&$3.2 \pm 0.6 $ & 10 & 12& 11d & $204\pm 21 $& 9.7 & 535 \\ 45 Her &151525 & 5.2 &B9p &$9380\pm 310$ & $ 2.18\pm 0.13$&$4.7 \pm 1.2 $ & 11 & 14 & 2d3m & $146 \pm 38$& 3.8 & 231 \\ &171586 & 6.4&A2pvar &$8760\pm 290$ & $ 1.37\pm 0.10$&$2.1 \pm 0.5 $ & 10 & 5 & 5d & $375\pm 56 $& 6.6 & 868 \\ &171782 & 7.8 &A0p &$9660\pm 315$ & $ 1.76\pm 0.30$&$2.7 \pm 1.3 $ & 10 & 6 & 2d1m & $333 \pm 78$& 4.2 & 584 \\ 19 Lyr &179527 & 5.9 &B9sp &$10370\pm 340$ & $ 2.63\pm 0.16$&$6.4 \pm 1.9 $ & 11 & 11& 8d2m& $156 \pm 46$& 3.4 & 211 \\ 4 Cyg& 183056 & 5.1 &B9sp &$11710\pm 385$ & $ 2.69\pm 0.11$&$5.3 \pm 1.2 $ & 12 & 13& 13d & $290\pm 42 $& 6.9 & 680 \\ &204411 & 5.3 &A6pe &$8750\pm 290$ & $ 1.97\pm 0.07$&$4.2 \pm 0.8 $ & 9 & 12& 12d & $88 \pm 14 $& 6.0 & 198 \\ $\kappa$ Psc &220825 &4.9&A0p&$9450\pm 310$ & $ 1.40\pm 0.05$&$1.9 \pm 0.3 $ & 10 & 12& 12d & $312\pm 25 $& 12.8& 865 \\ \hline \end{tabular} \end{table*} \subsection{Physical properties of the sample} For each of the sample stars, we have determined effective temperature $T_{\rm eff}$, luminosity $L$ and radius $R$, to allow us to identify the appropriate line mask for Least-Squares Deconvolution (Sect. 2.5) and for determination of the rotational axis inclination for the dipole magnetic field model (Sect. 2.7). Effective temperatures of stars of our sample were derived using Geneva and Str\"omgren photometry (obtained from the General Catalogue of Photometric Data (GCPD); Mermilliod et al. 1997) using the calibrations of Hauck \& North (1982) and Moon \& Dworetsky (1985). Effective temperatures reported in Table 1 are the average of the two estimates when both were available, or the single one which could be derived when only one photometric set was available. We have assumed for $T_{\rm eff}$ an uncertainty of the order of 3\% for propagation of uncertainties in all calculations using the effective temperature. Luminosity was inferred using the GCPD-reported visual magnitude, the Hipparcos parallax and the bolometric correction relations of Balona (1994). Radius was then inferred directly from the luminosity and temperature via the Stefan-Boltzmann equation for a uniform spherical star. The inferred values of $T_{\rm eff}$, $\log L/L_\odot$ and $R/R_\odot$ are reported in Table 1. For the 20 stars which we have in common with the study of Kochukhov \& Bagnulo (2006), these values are all in good agreement. As pointed out by Landstreet et al. (2007), the assumed uncertainties of Kochukhov \& Bagnulo (2006), which are comparable to our own, are probably somewhat underestimated. However, as our fundamental parameters are not to be used for detailed evolutionary studies, we consider them to be sufficient for this study. \subsection{Least-Squares Deconvolution and magnetic field detection} The primary aim of our study is to detect line circular polarisation (a ``Stokes $V$ Zeeman signature'') which is characteristic of the longitudinal Zeeman effect produced by the presence of a magnetic field in the stellar photosphere. For this we used the Least-Squares Deconvolution (LSD) procedure, first used by Donati et al. (1997) to study the magnetic fields of active late-type stars and by Wade et al. (2000 a,b) for Ap stars. This method enables the ``averaging'' of several hundred (and possibly several thousand in some stars) lines and thus to obtain Stokes $I$ and Stokes $V$ profiles with greatly improved S/N. LSD provides a single quantitative criterion for the detection of Stokes $V$ Zeeman signatures: we perform a statistical test in which the reduced $\chi^2$ statistic is computed for the Stokes $V$ profile, both inside and outside the spectral line (Donati et al. 1997). The statistics are then converted into detection probabilities, which are assessed to determine if we have a definite detection (dd, false alarm probability smaller than $10^{-5}$), a marginal detection (md, false alarm probability greater than $10^{-5}$ and smaller than $10^{-3}$), or no detection at all (nd). A diagnostic null spectrum (called $N$ in the following) is also obtained using the same subexposures obtained for Stokes $V$, but by pair processing those corresponding to identical azimuths of the quarter-wave plate. By checking that a signal is detected only in $V$ and not in $N$, and that any detected signature is located within the line profile velocity interval, we can distinguish between real magnetic signatures and (infrequent) spurious signatures. In addition, using the full resolved Stokes $V$ profile enables the detection of the magnetic field, even if the integrated line-of-sight component is very weak, or even null. \subsection {LSD masks} LSD is a cross-correlation method which requires comparison of our observed spectra with synthetic line masks (Donati et al. 1997, Shorlin et al. 2002). To obtain the most realistic masks, we used spectral line lists from the Vienna Atomic Line Database (VALD; Piskunov et al 1995; Ryabchikova et al. 1997; Kupka et al. 1999). To take into account the chemical peculiarities of Ap/Bp stars, we employed an abundance table in which the abundance of metals (Al, Si, S, Ti, V, Mn, Fe, Co, Ni, Zn, Sr, Y, Zr, Ba, La, Ce, Pr, Nd, Eu, Gd, Dy) is 10x solar, except for Cr which was increased to 100x solar (e.g. Shorlin et al. 2002). Masks were then compiled for effective temperature ranging from 7000-13000~K, with $\log g=4.0$, a microturbulence of 2~km/s, and including all metal lines with a central depth greater than 10\% of the continuum. We also compiled a series of masks assuming solar abundances. We have found in other studies that the magnetic field measurements are not very sensitive to the mask temperature within a couple of thousand K (Wade et al., in preparation). We therefore computed masks spaced every 1000~K. Least-Squares Deconvolution was performed for several temperatures and in some cases the most significant magnetic field detection was obtained for a temperature somewhat hotter than that given by the photometric data. This occured several times for the hottest sample stars - this can be seen in Table 1. In the case of the cool Ap star EP UMa, a solar abundance mask gave a better result (better detection of magnetic field and smaller error bars on $B_\ell$) than Ap abundances as described above. These discrepancies probably result from differences between the true chemical peculiarities of individual stars and those assumed in the line masks. For the discrepant hot stars, we computed the longitudinal magnetic field for a mask temperature between that corresponding to the best detection and the derived effective temperature. For EP UMa, we used solar abundance masks in our analysis. The number of lines used in the LSD ranged from 1500 to 3000, and is anticorrelated with the temperature. \subsection{Longitudinal magnetic field} The longitudinal magnetic field was inferred from each of the Stokes $I$ and $V$ profile sets, using the first-order moment method. According to this method, the longitudinal field $B_\ell$ (in G) is calculated from the Stokes $I$ and $V$ profiles in velocity units as: \begin{equation} B_{\ell} = -2.14 \times 10^{11} \frac{\int_{}^{}vV(v)dv}{\lambda gc\int_{}^{}[I_c - I(v)]dv}, \end{equation} \noindent (Rees \& Semel 1979; Donati et al. 1997; Wade et al. 2000b) where $\lambda$, in nm, is the mean wavelength of the LSD profile, $c$ is the velocity of light (in the same units as $v$), and $g$ is the mean value of the Land\'e factors of all lines used to construct the LSD profile. Integration ranges used for evaluation of Eq. (1) were computed automatically, beginning and ending 15~km/s before/after the location in the line wings at which the residual flux was equal to 85\% of the continuum flux. The accuracy of this technique for determining high-precision longitudinal field measurements has been clearly demonstrated by Wade et al. (2000b), Donati et al. (2001) and Shorlin et al. (2002). The resultant longitudinal magnetic field measurements, which are reported in Table~3, are remarkably precise (this Table is only available on line). The 282 measurements, with a median $1\sigma$ uncertainty of 40 G, represent the largest compilation of high-precision stellar magnetic field measurements ever published. \subsection{Modeling the longitudinal field variation} To characterise the dipole components of the magnetic fields of our sample stars, we use the oblique rotator model (ORM, Stibbs 1950) as formulated by Preston (1967). This model provides a good first approximation of the large scale magnetic field of Ap stars (e.g. Landstreet 1988). Because of the weakness of the longitudinal magnetic field observed for the stars of our sample, we do not expect to be able to detect departures from a global dipolar configuration. To begin, we added to our high-precision data set additional good-quality published magnetic field measurements collected by Bychkov et al. (2003). Details of these collected measurements were kindly provided by Dr. Victor Bychkov. Then, we searched the literature for rotational periods for each of our sample stars. For many stars, published rotational periods were available which provided an acceptable folded phase variation $B_\ell(\phi)$ of the magnetic measurements. However, for some stars, the published rotational period or periods did not provide an acceptable folded magnetic field variation, and for others, no published period was available. For these latter stars, we used a modified Lomb-Scargle technique to attempt to infer the rotational period, both directly from the longitudinal field measurements, as well as from the variations of the LSD Stokes $I$ and $V$ profiles. The period searches of LSD profiles were performed by treating each pixel in the Stokes $I$ and $V$ profiles as an independent timeseries (similar to the technique described by Adelman et al., 2002). Individual periodograms were subsequently weighted according to their amplitude of variation and averaged to characterise variability of the whole LSD profile. Acceptable periods were identified by establishing the 99\% confidence threshold, and candidate periods were evaluated by phasing the LSD profiles and longitudinal field measurements. Results of the period searches for individual stars are provided in their appropriate subsections in Sect. 3. Ultimately, acceptable rotational periods were obtained for 24 stars, and these periods are reported in Table 2. We also searched the literature for published values of the projected rotational velocity ($v\sin i$) of each star, which we compared to the value measured from the LSD Stokes $I$ profile by fitting rotationally-broadened synthetic profiles. Sometimes significant discrepancies were found between our values of $v\sin i$ and those reported in the literature. These discrepacies are discussed in Sect. 3, and the adopted rotational velocities (generally those obtained from the LSD profiles) are shown in Table 2. Each phased longitudinal field variation $B_\ell(\phi)$ was then fit using a 1$^{\rm st}$ order sine function: \begin{equation} B_\ell(\phi) = B_0 + B_1\sin {2\pi(\phi+\phi_0)}. \end{equation} The phased and fit longitudinal field variations are shown in Fig. 1. The reduced $\chi^2$ of this fit ($\chi^2_2$), along with those of linear fits through $B_\ell=0$ (the ``null field'' model, $\chi^2_0$ ) and through the weighted mean of the measurements (the ``constant field'' model, $\chi^2_1$), of each star, are reported in Table 2. $\chi^2_{\rm lim}$, the $2\sigma$ upper limit for admissible models, computed according to Press et al. (1992), is reported in Table 2 as well . A comparison of these reduced $\chi^2$ values with each other allows us to evaluate the significance of the detection of the longitudinal magnetic field and its variability. Although variability of the longitudinal field cannot be established for a few stars ($\chi^2_1<\chi^2_{\rm lim}$), the only star for which the longitudinal field is not detected with more than 2$\sigma$ confidence is HD 96707 (for which $\chi^2_0<\chi^2_{\rm lim}$). \begin{table*}[ht] \caption{Results of the magnetic field modeling. The contents of the columns are described in sect. 2.7. The uncertainties associated with the derived dipole parameters $i, \beta$ and $B_{\rm d}$ correspond to 2$\sigma$.} \begin{tabular}{lcc|cccc|cccccc} \\ \hline star & Period & $v\sin i$& $\chi^2_0$ & $\chi^2_1$ &$\chi^2_2$ &$\chi^2_{\rm lim}$& $i$ &$\beta$& $B_{\rm d}$& $B_{\rm d}^{\rm min}$& $B_{\rm d}^{\rm max}$\\ & (d) & (km/s) & & & & & ($\degr$) & ($\degr$) & ($10^3$ G) & ($10^3$ G) & ($10^3$ G)\\ \hline 8441 & 69.2 & 2 & 13.71 & 6.69 & 1.75 &3.35 &$ 49\pm33$ & $ 73_{- 61}^{+ 17}$ & 0.683 & 0.415 & 2.931 \\ 10221 & 3.15459 & 24 & 5.90 & 1.81 & 1.56 &2.71 &$ 27\pm11$ & $ 42_{- 41}^{+ 38}$ & 0.375 & 0.195 & 1.202 \\ 15089 & 1.74033 & 48 & 57.83 & 56.37 & 1.83 &2.14 &$ 45\pm11$ & $ 80_{- 12}^{+ 7}$ & 2.031 & 1.560 & 2.999 \\ 15144 & 2.99787 & 13 & 1305.83 & 4.57 & 0.04 &2.71 &$ 24\pm 8$ & $ 9_{- 3}^{+ 6}$ & 2.100 & 2.007 & 2.281 \\ 27309 & 1.568884 & 57 & 157.42 & 2.43 & 1.74 &2.63 &$ 53\pm17$ & $ 5_{- 5}^{+ 11}$ & 3.673 & 2.325 & 8.022 \\ 32549 & 4.6393 & 47 & 7.19 & 8.56 & 1.75 &2.75 &$ 65\pm27$ & $ 77_{- 74}^{+ 11}$ & 0.546 & 0.312 & 28.176 \\ 32650 & 2.7347 & 30 & 4.48 & 1.92 & 1.19 &1.72 &$ 37\pm12$ & $ 45_{- 30}^{+ 31}$ & 0.229 & 0.153 & 0.477 \\ 39317 & 2.6541 & 45 & 4.97 & 1.84 & 2.24 &3.18 &$ 36\pm12$ & $ 20_{- 20}^{+ 69}$ & 0.560 & 0.113 & 2.252 \\ 43819 & 15.02 & 10 & 135.66 & 59.32 & 3.45 &4.18 &$ 63\pm66$ & $ 42_{- 42}^{+ 47}$ & 2.626 & 2.488 & 78.367 \\ 68351 & 4.16 & 33 & 7.85 & 2.17 & 1.38 &2.00 &$ 28\pm37$ & $ 46_{- 41}^{+ 36}$ & 0.649 & 0.437 & 71.486 \\ 72968 & 5.6525 & 16 & 360.16 & 4.51 & 1.22 &2.03 &$ 61\pm18$ & $ 5_{- 4}^{+ 7}$ & 2.388 & 1.451 & 6.702 \\ 90569 & 1.04404 & 13 & 170.38 & 36.08 & 1.93 &3.07 &$ 9\pm 4$ & $ 81_{- 7}^{+ 5}$ & 5.157 & 2.946 & 11.284 \\ 94427 & 1.9625 & 8 & 37.58 & 46.61 & 2.12 &3.72 &$ 8\pm 4$ & $ 89_{- 4}^{+ 1}$ & 8.957 & 3.806 & 25.519 \\ 96707 & 3.515 & 37 & 0.82 & 0.91 & 0.77 &1.21 &$ 53\pm16$ & $ 90_{- 90}^{+ 0}$ & 0.100 & 0.0 & 0.492 \\ 103498 & 15.830 & 13 & 25.43 & 29.71 & 6.55 &7.44 &$ 75\pm68$ & $ 80_{- 11}^{+ 10}$ & 0.600 & 0.572 & 6.751 \\ 108945 & 2.01011 & 65 & 5.31 & 6.00 & 2.10 &2.83 &$ 57\pm18$ & $ 85_{- 61}^{+ 3}$ & 0.735 & 0.333 & 1.509 \\ 148112 & 3.04296 & 44.5 & 40.67 & 0.93 & 0.95 &1.73 &$ 56\pm16$ & $ 3_{- 3}^{+ 11}$ & 1.042 & 0.579 & 2.370 \\ 151525 & 4.1164 & 35 & 2.74 & 2.86 & 0.97 &1.70 &$ 37\pm19$ & $ 78_{- 43}^{+ 11}$ & 0.545 & 0.208 & 1.927 \\ 171586 & 2.1308 & 37 & 10.86 & 8.27 & 2.59 &6.60 &$ 48\pm19$ & $ 46_{- 40}^{+ 40}$ & 1.422 & 0.716 & 4.413 \\ 171782 & 4.4674 & 24 & 6.88 & 1.12 & 1.79 &3.79 &$ 51\pm51$ & $ 5_{- 5}^{+ 85}$ & 1.651 & 0.213 & 22257.276 \\ 179527 & 7.098 & 33 & 6.57 & 7.70 & 0.30 &1.30 &$ 74\pm32$ & $ 81_{- 34}^{+ 8}$ & 0.522 & 0.409 & 1.233 \\ 183056 & 2.9919 & 26 & 25.02 & 28.74 & 1.59 &2.32 &$ 74\pm 8$ & $ 49_{- 32}^{+ 37}$ & 1.558 & 1.172 & 3.938 \\ 204411 & 4.8456 & 5.4& 20.75 & 5.43 & 0.57 &1.46 &$ 7\pm 5$ & $ 81_{- 12}^{+ 7}$ & 0.968 & 0.416 & 4.509 \\ 220825 & 1.42539 & 38 & 78.54 & 90.05 & 2.29 &3.18 &$ 35\pm54$ & $ 83_{- 80}^{+ 7}$ & 1.957 & 1.141 & 21.045 \\ \\ \hline \end{tabular} \end{table*} For a tilted, centred magnetic dipole, the surface polar field strength $B_{\rm d}$ is derived from the variation of the longitudinal magnetic field $B_\ell$ with rotational phase $\phi$ using Preston's (1967) well-known relation: \begin{equation} B_d = B_\ell^{\rm max} \biggl({{15 + u}\over{20(3-u)}} (\cos\beta\cos i + \sin\beta\sin i)\biggr)^{-1}, \end{equation} \noindent where $B_\ell^{\rm max}=|B_0|+B_1$ and $u$ denotes the limb darkening parameter (equal to approximately $u=0.5$ for our sample). The rotational axis inclination and obliquity angles $i$ and $\beta$ are related by \begin{equation} \tan\beta={{1-r}\over{1+r}} \cot i, \end{equation} \noindent where $r=(|B_0|-B_1)/(|B_0|+B_1)$. We have determined the inclination $i$ for each of our stars assuming rigid rotation, and computing: \begin{equation} \sin i = {{P_{\rm rot} v\sin i}\over {50.6 R}}, \end{equation} \noindent where $P_{\rm rot}$ is the adopted stellar rotational period in days, $v\sin i$ is the adopted projected rotational velocity in km/s, and $R$ is the computed stellar radius in solar units. The magnetic obliquity $\beta$ was then inferred from Eq. (4) and the polar strength of the dipole $B_{\rm d}$ from Eq. (3). Uncertainties associated with all parameters (at 2$\sigma$, including a max$(2~{\rm km/s}, 10\%)$ uncertainty on $v\sin i$) were propagated through the calculations of $i$, $\beta$ and $B_{\rm d}$. The resultant dipole magnetic field models are reported in Table 2. An important and interesting consequence of Eq. (3), independent of all parameters except limb-darkening, is: \begin{equation} B_{\rm d} \geq {{20(3-u)}\over{15+u}} B_\ell^{\rm max}, \end{equation} \noindent which, for typical limb-darkening $u$, yields: \begin{equation} B_{\rm d} \gtrsim 3.3 B_\ell^{\rm max}. \end{equation} Therefore, exclusive of the model geometry, a lower limit on the surface dipole component of the magnetic field is obtained directly from the maximum measured value of the longitudinal field. The maximum measured longitudinal field is reported in column 11 of Table 1, and the inferred lower limit ($2\sigma$) of $B_{\rm d}$, $B_{\rm d}^{\rm min, 3.3}$ is reported in column 13 of Table 1. It is clear from these data that $B_{\rm d}$ for most of our sample is larger than a few hundred G at $2\sigma$. | We have investigated a sample of 28 well-known spectroscopically-identified magnetic Ap/Bp stars, obtaining 282 new Stokes $V$ Zeeman signatures and longitudinal magnetic field measurements using the MuSiCoS spectropolarimeter. Magnetic field is detected in all sample stars, and the inferred longitudinal fields are significantly greater than some tens of G. To characterise the surface magnetic field intensities of the sample, we modeled the longitudinal field data to infer the intensity of the dipolar field component. The distribution of derived dipole strengths for these stars exhibits a plateau at about 1 kG, falling off to larger and smaller field strengths. Remarkably, in this sample of stars selected for their presumably weak magnetic fields, we find only 2 stars for which the derived dipole strength is weaker than 300~G. Interpreting this ``magnetic threshold'' as a critical value necessary for the stability of large-scale magnetic fields leads to a natural explanation of the small fraction of intermediate-mass magnetic stars. It may also explain the near-absence of magnetic fields in more massive B and O-type stars. | 7 | 10 | 0710.1554 |
0710 | 0710.3883_arXiv.txt | Current observational cosmology allows us to test fundamental physics with a continuously improving precision. To evaluate and understand these data, theoretical models about the evolution of our universe are crucial. One promising paradigm receiving growing experimental support is cosmological inflation \cite{original_inflation}. Inflation postulates a period of exponential expansion of the universe driven by scalar fields slowly rolling in an almost flat potential. This enormous growth stretches quantum fluctuations present in the early universe to currently observable astrophysical scales. The imprints of such a process can be found, for example, in the cosmic microwave background and the large scale structure of the universe \cite{WMAP,Tegmark, Sanchez}. In recent years much effort has focused on the realization of inflation within string theory \cite{Cosmo_rev,Hertzberg:2007ke}. In string theory there are potentially many scalar fields which could drive inflation and hence different opportunities to model inflation. Specific scenarios include realizations of K\"ahler moduli inflation~\cite{Banks:1995dp,Conlon:2005jm}, racetrack models \cite{BlancoPillado:2006he} and the most intensively studied possibility of D-brane inflation \cite{Dvali:1998pa,Cosmo_rev}. However, as of today, it remains challenging to establish explicit scenarios in a controlled compactification without employing extreme fine-tuning to obtain a sufficiently flat potential \cite{BDKMcAS}. Having found a realization of inflation reproducing the current cosmological observables, it is important to establish which models can incorporate possible future observations. For example, as argued in refs.~\cite{BMcA}, many string scenarios do not allow for a high ratio $r$ of gravitational waves produced in the early universe. The current experimental bound on gravitational waves is $r<0.3$ \cite{Tegmark}, but future experiments, including Planck, BICEP and Spider \cite{Efstathiou:2006ak}, might allow the observation of $r$ with a precision down to $r >0.01$. It is thus desirable to study string embeddings of inflation which can incorporate an $r$ observable in these experiments. Recent attempts to do that can be found, for example, in refs.~\cite{Olsson:2007he,Krause:2007jr,Becker:2007ui,Kobayashi:2007hm, KSS}. A possible scenario able to incorporate primordial gravitational waves was suggested by Dimopoulos, Kachru, McGreevy and Wacker \cite{DKMW}. The authors argue for an embedding of multi-field inflation with a large number $N$ of axion fields. Such models use an assistance effect studied in refs.~\cite{LMS,KO} ensuring that the small fraction $1/N$ controls the flatness of the potential. Indeed, generic compactifications of type II string theory on a six-dimensional manifold can admit $10^4$, or more, axions from the NS-NS B-field and the R-R form fields. An appropriate subset of $N$ such axions were proposed to drive inflation in ref.~\cite{DKMW} and the authors termed these scenarios $N$-flation. For a sufficiently large $N$ the scenarios can be interpreted as a realization of natural inflation \cite{Freese:1990rb,Kim:2004rp}. If the inflatons can produce the desired amount of e-foldings already in the quadratic regime of the potential the models admit chaotic inflation \cite{Linde:1983gd}. This implies that $N$-flation can yield a possibly observable signature of gravitational waves with $r<0.14$ and thus distinguishes it from most other string realizations of inflation. The aim of this paper is to study a specific realization of $N$-flation in type IIB string theory. The inflating axions correspond to the zero-modes of the R-R forms in the compactification to four space-time dimensions. In compactifications preserving $\cN=2$ supersymmetry the R-R axions sit in the same supermultiplets as the K\"ahler structure moduli parametrizing the volumes of two-dimensional cycles in the internal space. Manifolds with a large number of non-trivial two-cycles are thus candidate backgrounds for $N$-flation. As will be shown, the density perturbations and slow roll parameters depend on the volume of the compact space and it has to be ensured that they do not become large with increasing $N$. We will thus argue that explicit examples always involve compact manifolds which admit many very small or vanishing cycles. In the presence of small cycles stringy effects become important and significantly alter the structure of the four-dimensional effective theory. In order to analyze these contributions we will concentrate on axions arising from the R-R two-form evaluated on vanishing two-cycles of the compact geometry. The standard example of a vanishing two-cycle is the resolved conifold \cite{Candelas:1989js}. In this case a conical singularity is resolved by a two-sphere supported by a geometric volume or an NS-NS B-field. If this $S^2$ becomes smaller than the square string length, world-sheets will start to wrap and significantly contribute to the metric of the R-R axions. Fortunately, in $\cN=2$ compactifications these corrections can be computed for the conifold and many other Calabi-Yau geometries \cite{Hori:2003ic,Hosono:1993qy} allowing the evaluation of the metric of the axions. We will illustrate this general fact on a toy model with $N$ conifold singularities. For such examples it can be shown that the kinetic terms of the axions can be close to the Planck scale which is crucial to obtain inflation \cite{BDFG}. In addition to the fundamental strings also D-branes can become relevant in geometries with vanishing cycles \cite{Strominger:1995cz}. In particular, D1 instantons can wrap the small cycles and correct the metric of the R-R axions. Such contributions appear with the exponential of the D1 instanton action which depends on the R-R two-form axions themselves. In contrast to the string world-sheet action the D-instanton action also contains a factor of the inverse string coupling. This implies that for small string coupling and finite volume or B-field of the vanishing cycles the D1 instanton corrections are subleading in the axion metric. However, correction due to D-branes will be the leading contributions in the scalar potential and can induce the desired potential for the R-R axions. To study the scalar potential for the axions and non-axionic moduli we will focus on $\cN=1$ orientifold compactifications of type IIB string theory. Such compactifications have been studied intensively in the last years \cite{review_flux}. It was shown that $\cN=1$ potentials can be induced by background fluxes, D-brane instantons or gaugino condensates on space-time filling D-branes. The desired axion potentials can arise through non-trivial superpotentials from either of the three sources \cite{review_flux}. We will briefly discuss their properties in various orientifold compactifications: D1 superpotentials in type I~\cite{Witten2}, D1 dependences through the determinants in the D3 instanton superpotentials \cite{Witten3,Ganor,TG}, and gaugino condensates on space-time filling D5 branes \cite{review_flux,Vafa:2000wi,Cachazo:2001jy,Dijkgraaf:2002dh}. Remarkably, explicit computations of the non-perturbative superpotentials can often be performed in a dual flux picture where geometry dictates the form of the corrections \cite{Vafa:2000wi,Cachazo:2001jy,Dijkgraaf:2002dh,Heckman:2007ub,ABK}. In the final part of this work we will study the effective theory of orientifold compactifications with O3 and O7 planes in more detail. Using earlier results \cite{GL1,GL2} we argue that the $\cN=2$ world-sheet corrections are inherited by the $\cN=1$ orientifold theory. In particular, they correct the $\cN=1$ K\"ahler potential and complex coordinates in a calculable way. This general fact can be applied to a simplistic compact toy model with $2N$ conifolds pairwise identified under the orientifold projection. Including a flux and D-instanton superpotential, we make first steps in establishing an effective theory with a large number of light axions and all other moduli stabilized. An explicit numerical evaluation indicates the presence of a non-supersymmetric axion valley \cite{Kallosh}. This ensures the desired mass hierarchy between the axions and their non-axionic partners. The paper is organized as follows. In section \ref{rev_N} we review the $N$-flation scenario of \cite{DKMW} and discuss some of its cosmological implications. We recall that the kinetic terms of the axions, set by the axion decay constants, have to be large in order to obtain inflation. A discussion of axion decay constants in Type IIB string theory is presented in section~\ref{axion_decay}. The four-dimensional effective Lagrangian and the general form of the axion decay constants are studied in section \ref{four-dim_Lagr} supplemented with appendix \ref{gen_axion_const}. In section \ref{large_vol} we argue that in compactifications with all cycles larger than string length, the axion decay constants typically become very small with an increasing number of axions. This implies that only compactification manifolds with small or vanishing cycles are candidate backgrounds to obtain $N$-flation. The quantum corrected axion decay constants for the resolved conifold and geometries with $N$ resolved singularities are discussed in sections \ref{res_cone} and \ref{multiax}. Non-perturbative D-brane effects can induce a scalar potential through a non-vanishing superpotential as discussed in section \ref{axion_potential}. In section \ref{D1_corrections} it is shown that such a superpotential can arise from D1 instanton corrections, while section~\ref{D5_gaugino} discusses superpotentials originating from gaugino condensates on D5 branes. In the final part of this work, section \ref{N=1_ori}, the embedding of $N$-flation into a concrete $\cN=1$ orientifold compactification is addressed. The general form of the $\cN=1$ data including the inherited $\cN=2$ perturbative and non-perturbative string world-sheet corrections is presented in section~\ref{four-dim_N=1} and appendix~\ref{N=1_recall}. In section~\ref{N_conifolds}, a toy model with $N$ conifold pairs is used to illustrate that the string world-sheet corrections in the K\"ahler potential and the presence of a non-perturbative superpotential can ensure an effective theory with light axions. This indicates the possibility of $N$-flation in string theory. | In this paper we discussed the possibility of a string theoretical embedding of cosmological inflation driven by a large number of axions. Such scenarios use the fact that in the dimensional reduction of string theory to four space-time dimensions a vast number of scalar fields arise in the low energy theory. In particular, we considered type IIB string theory on a compact Calabi-Yau manifold with many non-trivial two-cycles. To each of these two-cycles an axion from the R-R two-form and an axion from the R-R four-form can be associated. The effective theory for the axionic fields will sensitively depend on the values of the geometric moduli, i.e.~the volume of the two-cycles, as well as the NS-NS B-field. We have argued that a possible realization of axion inflation might only exist in special corners in the landscape of vacua. In these regimes various stringy effects become relevant and have to be included. In this work, we have made first steps one region in the $\cN=1$ landscape, where axion inflation might be realized. In recent years, $\cN=1$ type IIB compactifications with all volumes stabilized at scales much larger than string scale have been investigated intensively \cite{review_flux}. For vacua in this regime of the field space, stringy corrections, such as wrapping stings and D-branes, play a sub-leading role in the derivation of the kinetic terms of the axions. Non-perturbative D-brane effects can induce a potential exponentially suppressed by the large volume of the cycles times the inverse of a small string coupling. We have argued that such scenarios are not suitable for axion inflation. They naturally admit very small axion decay constants which become even smaller if the number of axions increases. Moreover, due to the strong exponential suppression, the scale of inflation is typically too low for semi-realistic scenarios. This conclusion lead us to the consideration of manifolds with small or vanishing cycles. In compactifications with cycles smaller than string scale the caveats of the large volume compactifications can be avoided. To obtain such geometries, we considered resolutions of singularities supported by a volume and a NS-NS B-field. The standard example of such a blown-up singularity is the resolved conifold. In the vicinity of small cycles, new corrections will become relevant and alter the effective theory. We have discussed a subset of such correction in $\cN=2$ compactifications of type IIB string theory on a Calabi-Yau manifold. Including the leading singular non-perturbative string world-sheet contributions, we have argued that the axion decay constants can take values close to the Planck scale also for scenarios with many axions. In $\cN=1$ compactifications further perturbative and non-perturbative corrections can become relevant and might alter the structure of the effective four-dimensional theory. However, in case these do not cancel the $\cN=2$ effects, large axion decay constants will remain accessible also in these less supersymmetric scenarios. In addition to being close to the Planck scale, $N$-flation also requires the axion decay constants to be independent of the axions themselves. Corrections depending on the R-R axions are believed to only arise from non-perturbative D-brane effects. If the small cycles remain to be of finite size and we work at sufficiently small string coupling, D-brane corrections are subleading in the axion decay constants, since the instanton action is larger by a factor of the inverse string coupling. This will also be the case for vacua in an $\cN=1$ supergravity theory. In a consistent analysis, the vacuum values of the $\cN=1$ moduli fields have to be inside an appropriate field range to ensure that D-brane instantons are subleading in the K\"ahler potential encoding the kinetic terms of the scalar fields. One of the remaining complications in realizing $N$-flation is to ensure that there indeed exists an effective theory for axions with appropriate masses during inflation. These masses have to be lower than the masses of their non-axionic partner and other bulk moduli fields, but still sufficiently large to match the observed density perturbations during inflation. Again, this forces us to work away from the large volume regime, where at least some of the moduli masses are suppressed by large instanton actions. A careful investigation of the effective potential for the moduli is necessary to study the vacua of the theory. In this work we briefly discussed effective $\cN=1$ potentials arising from D-instantons and gaugino condensates on space-time filling branes. We pointed out that for potentials induced by large rank gaugino condensates on D5 branes, the quadratic region of the axion potential can stretch over the entire accessible field range. In practice, D1 corrections arising through the pre-factors of the D3-instantons are of particular interest. If these are the leading axion-dependent contributions to the $\cN=1$ superpotential, a mass hierarchy can be ensured due to the suppression by both the D1 and D3 instanton action. To investigate the properties of the effective theory more explicitly, we discussed the embedding of axion inflation into $\cN=1$ Calabi-Yau orientifold compactifications with O3 and O7 planes. We showed that the $\cN=1$ characteristic data remain calculable even in the case that the non-perturbative $\cN=2$ string world-sheet corrections are included. Utilizing a flux superpotential together with a superpotential from non-perturbative D-brane effects a potential for all bulk moduli fields is generated. We illustrated that a theory of light axions could exist, if the axion dependence is suppressed in the superpotential and the non-axionic partners of the axions in the $\cN=1$ chiral multiplet is stabilized due to the string world-sheet corrections to the K\"ahler potential. In an optimistic scenario, one can hope that such an effective theory of a large number of relatively light axions from the R-R forms will survive also further perturbative corrections. Even though an explicit embedding of $N$-flation into string theory still remains to be constructed, the scenarios outlined and studied in this work might provide a promising route to achieve this goal. Likely, such an embedding will not solve intrinsic issues related to the fine-tuning of initial conditions in chaotic and natural inflation with many inflatons. However, it might provide a way to accommodate possible future observations of primordial gravitational waves in a string theoretic model. | 7 | 10 | 0710.3883 |
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0710 | 0710.5188_arXiv.txt | By-now photons are the unique universal messengers. Cosmological sources like far-away galaxies or quasars are well-known light-emitters. Here we demonstrate that the nonlinear electrodynamics (NLED) description of photon propagation through the weak background intergalactic magnetic fields modifies in a fundamental way the cosmological redshift that a direct computation within a specific cosmological model can abscribe to a distant source. Independently of the class of NLED Lagrangian, the effective redshift turns out to be $1 + \tilde{z} = (1 + z)~\Delta$, where $\Delta \equiv (1 + \Phi_e)/(1 + \Phi_o)$, with $\Phi \equiv {8}/{3} ({L_{FF}}/{L_F}) B^2$, being $L_F = {dL}/{dF}$, $L_{FF} = {d^2L }/{dF^2}$, the field $F\equiv F_{\alpha \beta} F^{\alpha \beta}$, and $B$ the magnetic field strength. Thus the effective redshift is always much higher then the standard redshift, but recovers such limit when the NLED correction $\Delta(\Phi_e, \Phi_o) \longrightarrow 1$. This result may provide a physical foundation for the current observation-inspired interpretation that the universe undergoes an accelerate expansion. However, under the situation analyzed here, for any NLED the actual (spatial) position of the light-emitting far-away source remains untouched. | 7 | 10 | 0710.5188 |
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0710 | 0710.2761_arXiv.txt | {(Ultra)Luminous Infrared Galaxies [(U)LIRGs] are much more numerous at cosmological distances than locally, and are likely the precursors of elliptical galaxies. Studies of the physical structure and kinematics of representative samples of these galaxies at low redshift are needed in order to understand the interrelated physical processes involved. Integral field spectroscopy (IFS) offers the possibility of performing such a detailed analysis.} {Our goal is to carry out IFS of 42 southern systems which are part of a representative sample of about 70 low redshift (z $\le$ 0.26) (U)LIRGs, covering different morphologies from spirals to mergers over the entire infrared luminosity range (L$_{IR}$=10$^{11}-10^{12.6}$L$_{\odot}$).} {The present study is based on optical IFS obtained with the VIMOS instrument on the VLT.} {The first results of the survey are presented here with the study of two galaxies representative of the two major morphological types observed in (U)LIRGs, interacting pairs and morphologically- regular, weakly-interacting spirals, respectively. We have found that IRAS F06076$-$2139 consists of two low-intermediate mass (0.15 and 0.4 m$_*$) galaxies with relative velocities of $\sim$ 550 km s$^{-1}$ and, therefore, it is unlikely that they will ever merge. The VIMOS IFS has also discovered the presence of a rapidly expanding and rotating ring of gas in the southern galaxy (G$_s$). This ring is interpreted as the result of a nearly central head-on passage of an intruder about 140 million years ago. The mass, location and relative velocity of the northern galaxy (G$_n$) rules out this galaxy as the intruder. IRASF 12115$-$4656 is a spiral for which we have found a mass of 1.2 m$_*$. The ionized gas shows all the kinematic characteristics of a massive (8.7 $\times$ 10$^{10}$M$_{\odot}$), fast rotating disk. The neutral gas traced by the NaI doublet shows distinct features not completely compatible with pure rotation. The neutral and ionized gas components are spatially and kinematically decoupled. The analysis presented here illustrates the full potential of IFS in two important aspects: (i) the study of the kinematics and ionization structure of complex interacting/merging systems, and (ii) the study of the kinematics of the different gas phases, neutral (cool) and ionized (warm), traced by the NaD and H$\alpha$ lines, respectively.} {} | Most of the observational efforts devoted in recent years to the study of galaxy formation and evolution have been focused on obtaining and analyzing large surveys of different characteristics such as, for instance, 2MASS (Skrutskie et al. 2006), SDSS (Adelman-McCarthy et al. 2006), 2DF (Colless et al. 2001), UDF (Beckwith et al. 2006), GOODS (Giavalisco et al. 2004), COSMOS (Capak et al. 2007), GEMS (Rix et al. 2004), etc. Thanks to these studies observational constrains on integrated properties such as the galaxy luminosity functions, galaxy (stellar) masses, sizes, spectral energy distributions, etc., as well as their evolution with redshift, have drastically improved (see, for instance, Rix et al. 2006 and references therein). However, the observational limits imposed by the HST and 10 meter class telescopes have been reached and a substantial new improvement using this methodology will probably have to wait until a new generation of telescopes (JWST, ELT, etc) comes into operation. In addition to those integrated properties, a comprehensive picture of galaxy formation should explain the internal structure of galaxies (i.e., internal kinematics, internal stellar population gradients, dust distribution, ionization structure, nuclear properties and effects, interaction with the IGM, etc). Such detailed studies are more demanding in terms of 'data quality per object' requiring high S/N two dimensional spectroscopic information with high angular and spectral resolution. Thanks to the advent and popularization of Integral Field Spectroscopic instruments it is now possible to obtain high quality two-dimensional spectroscopy (i.e. 3D data) over a wide wavelength range ($>$ 1000 $\AA$) in one spectral setting. Although most such studies have been based on local samples of galaxies (e.g., Mediavilla et al. 1997; Peletier et al. 1999, 2007; del Burgo et al. 2001; de Zeeuw et al. 2002 and references therein), recent work also includes high-z studies (e.g., Smith et al. 2004; Forster-Schreiber et al. 2006; Flores et al. 2006; Puech et al. 2007; Law et al. 2007). However, limitations in sensitivity and angular resolution make the studies of the internal structure of high-z galaxy populations via IFS extremely difficult, and the observation of comprehensive samples (i.e., not just individuals at the tip of the luminosity function) will require larger aperture telescopes, and higher angular resolution than those currently available. In the context of the high-z universe, the study of representative samples of some low-z galaxies populations are of particular importance. That is the case of the luminous and ultraluminous infrared galaxies (LIRGs, L$_{IR}$=L[8-1000$\mu$m]=10$^{11}$-10$^{12}$L$_{\odot}$, and ULIRGs, L$_{IR}$$>$10$^{12}-10^{13}$L$_{\odot}$, respectively). These objects are thought to be the local counterparts of cosmologically important galaxy populations at z $>$ 1 such as the submillimiter galaxies (SMG; e.g., Frayer et al. 2004), and the majority of the infrared sources found in recent {\it Spitzer} cosmological surveys (e.g., P\'erez-Gonz\'alez et al. 2005). ULIRGs and LIRGs have large amounts of gas and dust, and are undergoing an intense star formation in their (circum)nuclear regions (e.g. Scoville et al. 2000, Alonso-Herrero et al. 2006, and references therein). This starburst activity is believed to be their major energy reservoir, although AGN activity may also be present (Genzel et al. 1998). In many cases (especially for ULIRGs) these objects show clear signs of an on-going merging process producing moderate-mass (0.1$-$1 m$_*$) ellipticals as the end product of the merger (Dasyra et al. 2006, Genzel et al. 2001). Although we have a qualitative picture, details on how the physical and dynamical processes involved are interrelated and how they determine the internal structure of these cosmologically important objects are far from being well understood. We have already started a program aimed at studying the internal structure and kinematics of a representative sample of low-redshift LIRGs and ULIRGs using optical and near-IR Integral Field Spectroscopy. These objects are natural laboratories that allow to investigate in detail the internal structure of these galaxy populations in a high S/N, high spatial resolution regime, capturing their complex two dimensional structure. We are mainly analyzing rest-frame optical and near-IR spectral diagnostic features on linear scales of about 1 kpc, that will appear in the near- and mid-IR in the high-z galaxy populations to be studied with future instruments such as the NIRSpec and MIRI instruments on board the JWST. This program initially has been focused on the most luminous objects, the ULIRGs (e.g., Colina, Arribas \& Monreal-Ibero, 2005 and references therein; Garc\'{\i}a-Mar\'\i n 2007) and now is being extended towards lower luminosities, i.e. the LIRG luminosity range. The present paper is the first of a series of studies based on a representative sample of LIRGs observed with VLT-VIMOS. With $\sim$ 2000 spectra per object, this study involves a large analysis effort. In this first paper we describe the sample, the experimental set-up used for all the observations with VIMOS, as well as the reduction procedures and analysis methods. The analysis of two of the galaxies representative of the sample is shown to illustrate and emphasize the power of the survey in several areas, which after completion will include, among others, the following: (i) two dimensional kinematics of different phases (neutral and ionized) of the interstellar gas as a function of luminosity and morphology, (ii) two dimensional ionization structure, and the interplay between ionization and kinematics (i.e., the role of shocks), and (iii) star formation in the (circum)nuclear and external regions (candidate Tidal Dwarf Galaxies and/or precursors of globular clusters) in interacting/merging systems. Throughout the paper we will consider H$_{0}$=70 km~s~$^{-1}$Mpc$^{-1}$, $\Omega_{M}$= 0.7, $\Omega_{\Lambda}$= 0.3. | This paper presents the first results obtained from the integral field spectroscopic survey of local luminous and ultraluminous infrared galaxies being carried out with the VIMOS instrument on the VLT. This sample includes 42 local (U)LIRGs systems (65 galaxies) and it is part of a larger representative sample of about 70 systems (z $<$ 0.26) also observed from the northern hemisphere, which covers the different morphologies from spirals to mergers and spans the entire luminosity range L$_{IR}$=10$^{11}-10^{12.6}$L$_{\odot}$. The two galaxies analysed here (IRAS F06076$-$2139 and IRAS F12115$-$4656) have been selected as prototypes of the most common morphological types among (U)LIRGs, i.e. interacting pairs and regular/weakly-interacting spirals, while also covering a wide range in IR luminosity. A) The VIMOS data of IRAS F06076$-$2139 clearly illustrate the potential of IFS in disentangling the complex ionizing and kinematic structure detected in pairs or mergers. The main conclusions for this system are: a1) We have found that this system is formed by two independent low-intermediate mass galaxies, the G$_n$ (northern) and G$_s$(southern), of 0.4 m$_*$ and 0.15 m$_*$, respectively, and with a relative projected velocity of about 550 km s$^{-1}$. Therefore, it is unlikely that these galaxies will ever merge. G$_n$ is characterized by an AGN-like ionization, and large velocity dispersions ($\geq$ 100 km s$^{-1}$), while G$_s$ has HII characteristics and lower velocity dispersions (by at least a factor of two) than G$_n$. a2) The morphology and velocity field of the ionized gas in G$_s$ is consistent with the presence of an expanding, rotating ring of about 20 kpc (deprojected major axis) in size. This expanding ring is interpreted as the result of a head-on collision with an intruder that took place some 1.4 $\times$ 10$^{8}$ years ago. The characteristics of G$_n$ rule out this galaxy as the intruder. a3) A chain of H$\alpha$ emitting regions kinematically associated with G$_s$ at distances of about 12 kpc have been detected. These regions are identified as young, tidally-induced star-forming regions (some may be candidate tidal dwarf galaxies) formed as a consequence of a past interaction. B) The VIMOS data of IRAS F12115$-$4656 show the potential of IFS in the kinematic study of the different phases of the interstellar gas: the neutral gas traced by the NaI absorption doublet, and the ionized gas traced by the H$\alpha$ or [NII] emission lines. The main conclusions for this system are: b1) The neutral and ionized gas morphology show a different structure indicating that they have a different spatial distribution. The neutral gas closely follows the stellar distribution while the ionized gas traces mostly the location of HII regions in the circumnuclear and external regions of the galaxy. b2) The ionized gas shows typical kinematic characteristics of massive rotating disks, with a large amplitude (peak-to-peak of 520 km s$^{-1}$), a maximum of the velocity dispersion at the kinematic/light centre, and a major kinematic axis coincident with the stellar photometric axis. A dynamical mass of 8.7 $\times$ 10$^{10}$ M$_{\odot}$ is measured inside a region of 6 kpc radius. The neutral gas, although consistent on a large scale with a rotating system, shows clear departures from pure rotation and it is dynamically hotter than the ionized gas, being more affected by random motions. b3) Despite being dominated by rotation the ionized gas shows large-scale (10 kpc), low-amplitude (20-30 km s$^{-1}$) non-rotational motion, which spatially correlates with regions of increased velocity dispersion. The neutral gas also exhibits an even clearer spatial correlation between departures from rotation and high velocity dispersion values. | 7 | 10 | 0710.2761 |
0710 | 0710.5836_arXiv.txt | { {\it Aims.} Tidally locked rotation is a frequently applied assumption that helps to measure masses of invisible compact companions in close binaries. The calculations of synchronization times are affected by large uncertainties in particular for stars with radiative envelopes calling for observational constraints. We aim at verifying tidally locked rotation for the binary PG~0101$+$039, a subdwarf B star + white dwarf binary from its tiny (0.025\%) light variations measured with the MOST satellite (Randall et al. \cite{randall}).\\ {\it Methods.} Binary parameters were derived from the mass function, apparent rotation and surface gravity of PG~0101$+$039 assuming a canonical mass of 0.47\,M$_{\rm \odot}$ and tidally locked rotation. The light curve was then synthesised and was found to match the observed amplitude well.\\ {\it Results.} We verified that the light variations are due to ellipsoidal deformation and that tidal synchronization is established for PG~0101$+$039. We conclude that this assumption should hold for all sdB binaries with orbital periods of less than half a day. Hence the masses can be derived from systems too faint to measure tiny light variations. | } The masses of compact objects like white dwarfs, neutron stars and black holes are fundamental to astrophysics, but very difficult to measure. Close binary systems consisting of a visible primary and an invisible compact object are very useful to this end as the companion mass can be constrained from the radial velocity and the light curve of the primary if the primary mass and the orbital inclination are known. The latter can be measured in systems that are eclipsing or show tidally locked rotation. If the mass of the primary can be estimated, the companion mass can then be derived. The assumption of tidally locked rotation has often been used to determine the masses of neutron stars and black holes in known X-ray binaries (see Charles \& Coe \cite{charles} for a review). The same technique has recently been applied to KPD~1930$+$2752, a short period binary consisting of a subluminous B stars and a white dwarf (Geier et al. \cite{geier1})\footnote{The companion is so massive that the system mass may exceed the Chandrasekhar mass, making the system a viable Supernova Ia progenitor candidate in the double degenerate scenario.}. sdB stars are core helium burning stars with very thin hydrogen envelopes and masses around $0.5M_{\rm \odot}$ (Heber et al. \cite{heber1}). A large fraction of the sdB stars are members of short period binaries (Maxted et. al \cite{maxted2}; Napiwotzki et al. \cite{napiwotzki6}). For these systems common envelope ejection is the most probable formation channel (Han et al. \cite{han2}). The companions of sdBs in these systems are predominantly white dwarfs implying that the system has undergone two phases of common envelope ejection. In order to establish bound rotation for an sdB star the synchronization time scale has to be smaller than its evolutionary life time ($t_{\rm EHB}\approx10^{8}\,{\rm yrs}$). Two theoretical concepts to compute synchronization times have been developed by Zahn (\cite{zahn}) and Tassoul \& Tassoul (\cite{tassoul}), respectively. While the mechanism proposed by Zahn is not efficient enough to fully fit the observed levels of synchronization, the more efficient one of Tassoul \& Tassoul is matter of a controversy, given its free parameter dependence (see Claret et al. \cite{claret1}, \cite{claret2} and references therein). Especially in the case of hot stars with radiative envelopes, where tidal forces are less effective in synchronizing the stars, the predictions of the two theoretical models at hand can differ by orders of magnitude. All prior studies were undertaken to match observations of hot main sequence stars. Hot subdwarf stars have similar temperatures as B-type main sequence stars, but are much smaller and the internal structure of these helium core burning objects is different. In addition, the fraction of sdBs residing in close binary systems is among the highest known of all types of stars. Observations of hot subdwarfs could provide a new benchmark to study the yet unresolved problem of tidal dissipation in radiative stellar envelopes. Independent observational constraints are needed to prove or disprove synchronized rotation in hot subdwarf stars. Ellipsoidal variations can be used to verify synchronization of the stellar surface because the light variations would then have to occur at exactly half the orbital period. Two sdB + white dwarf binaries are known to show ellipsoidal variations at half of the orbital period (KPD\,0422+5421 Orosz \& Wade \cite{orosz}; KPD\,1930+2752, Geier et al. \cite{geier1}). However, both systems have short orbital periods of about $0.1\,{\rm d}$ and theory predicts synchronization times much smaller the evolutionary time scale. As the synchronization time strongly increases with increasing period, we expect an upper limit to the period to exist at which the assumption of tidally locked rotation breaks down. To this end it would be of utmost importance to find ellipsoidal variations in an sdB binary of longer period and to provide a stringent test for the theory of synchronization. Recently a suitable object has been found. PG~0101$+$039, an sdB+WD binary (P=0.567 d, Maxted et al. \cite{maxted2}) was discovered to show very weak luminosity variation at half the orbital period in a light curve in a 16.9 day long, almost uninterrupted light curve obtained with the MOST satellite (Randall et al. \cite{randall}). In order to verify that we indeed see ellipsoidal variations, we have to show that the observed light curve can be consistently modelled. Beforehand, we have to derive the complete set of system parameters. As the spectrum is single lined, the analysis of the radial velocity curve yields the mass function only. Complementing it with an estimate of the sdB mass and with measurements of the sdB's projected rotational velocity as well as its gravity allows to solve for all binary parameters and compute the light curve. | Tidally locked rotation in close binary systems has been assumed to measure masses of invisible compact companions, in particular in X-ray binaries. The synchronization time is very difficult to calculate for stars with radiative envelopes and plagued with large uncertainties. Therefore, observational constraints are of utmost importance. Ellipsoidal variations can be used to verify the assumption at least for the surface layers. We applied this technique to the sdB/WD binary PG~0101$+$039, for which a very weak luminosity variation at half the orbital period has been discovered in a light curve in a 16.9 day long, almost uninterrupted light curve obtained with the MOST satellite. From spectroscopy we measured the mass function, apparent rotation and surface gravity of PG~0101$+$039. Stellar evolution models suggest that the sdB mass is close to 0.47\,M$_{\rm \odot}$. Assuming tidally locked rotation, this information is sufficient to solve for all parameters of the binary system. The companion mass is found to be $M_{\rm WD}=0.72 \pm 0.10\,M_{\rm \odot}$, typical for a white dwarf. The light curve was then synthesised and was found to match the observed amplitude well. However, a problem with the phasing of the light curve to the radial velocity curve became apparent. Due to a six year difference between the MOST photometry and published radial velocities, the phase errors were far too large for any conclusion to be drawn. Therefore we added 16 radial velocities from three observatories. The statistical error of the period decreased sufficiently to enable proper phasing of the photometry. The synthesised light curve was found to be offset by 0.1 cycles from the observed one indicating that our systematic error estimate may be overly optimistic. Alternative explanations like supersynchronous rotation of the sdB that may cause the observed phase shift seem to be unlikely because a deviation of $10\%$ from equilibrium would require fast rotation of the sdB. In this case the inclination would be very low and the companion mass would rise dramatically. A simultaneous measurement of the radial velocity curve and a high precision light curve would be necessary to solve this problem since the theoretical understanding of angular momentum transfer in hot stars with radiative envelopes is still very limited. In conclusion, we found strong indication that the surface rotation of the sdB star PG~0101$+$039 is tidally locked to its orbit. The synchronization times for any given type of primary depend strongly on the orbital period (Zahn \cite{zahn}, Tassoul \& Tassoul \cite{tassoul}). Hence, other sdB stars in close binaries should also be synchronized if their orbital period is less than that of PG~0101$+$039 (P\,=\,$0.567\,{\rm d}$). Hence we conclude that tidally locked surface rotation is established in sdB binaries with orbital periods of less than half a day. Hence the assumption of tidally locked rotation can be safely applied to such systems, even if they are too faint to measure such extremely small light variations as observed here. | 7 | 10 | 0710.5836 |
0710 | 0710.1718_arXiv.txt | We derive accretion rate functions (ARFs) and kinetic luminosity functions (KLF) for jet-launching supermassive black holes. The accretion rate as well as the kinetic power of an active galaxy is estimated from the radio emission of the jet. For compact low-power jets, we use the core radio emission while the jet power of high-power radio-loud quasars is estimated using the extended low-frequency emission to avoid beaming effects. We find that at low luminosities the ARF derived from the radio emission is in agreement with the measured bolometric luminosity function (BLF) of AGN, i.e., all low-luminosity AGN launch strong jets. We present a simple model, inspired by the analogy between X-ray binaries and AGN, that can reproduce both the measured ARF of jet-emitting sources as well as the BLF. The model suggests that the break in power law slope of the BLF is due to the inefficient accretion of strongly sub-Eddington sources. As our accretion measure is based on the jet power it also allows us to calculate the KLF and therefore the total kinetic power injected by jets into the ambient medium. We compare this with the kinetic power output from SNRs and XRBs, and determine its cosmological evolution. | \label{s:intro} There is increasing support for the idea that feedback from active galactic nuclei (AGN) plays an important role for galaxy formation \citep[e.g.,][]{CowieBinney1977,BinneyTabor1995,SilkRees1998,ChurazovKaiser2001,diMatteoSpringelHernquist2005}. This feedback is also invoked to explain the M-$\sigma$ relation of supermassive black holes \citep[e.g.,][]{HaehneltNatarajanRees1998,King2003,RobertsonHernquistCox2006,FabianCelottiErlund2006}. It is not yet clear whether kinetic or radiative feedback is dominant and how efficient each is at a given accretion rate. In this paper we will exploit the analogy between X-ray binaries (XRBs) and AGN to obtain information about the efficiency of the different feedback processes and calculate the total power available for feedback at a given redshift. The central engines of XRBs and AGN seem to be similar \citep[e.g.,][]{MirabelRodriguez1998,Meier2001} and recently relations have been found which scale spectral and variability properties from one class to the other \citep{MerloniHeinzdiMatteo2003,FalckeKoerdingMarkoff2004,KoerdingJesterFender2006,McHardyKoerding2006}. In the nearby universe ($z\la 0.2$), there are very few high-luminosity quasars like those that exist at high redshifts. These bright high redshift quasars are likely to have a strong effect on the evolution of the galaxy luminosity function. However, there are several XRBs that go through transient phases of very high accretion rates. These ``very-high state'' objects may be better templates for the bright quasars at a redshift of two than any nearby AGN. Thus, we will use our knowledge of XRBs and their states to obtain information about the kinetic and radiative properties of AGN. For XRBs one can observe a full outburst cycle, in which the accretion rate increases from very low ($10^{-7} \leq \frac{\dot{M}}{\dot{M}_{Edd}} \leq 10^{-5}$) to near the Eddington limit and then decays again. At low accretion rates, the source is generally found in the \emph{hard state}, which is characterized by a hard power law in the X-ray spectrum. The hard X-ray emission is usually accompanied by radio emission associated with a compact jet \citep{Fender2001} which sometimes can be directly imaged \citep[e.g.,]{StirlingSpencer2001}. The accretion flow in the hard state is likely to be radiatively inefficient \citep*[e.g.,][]{EsinMcClintockNarayan1997,KoerdingFenderMigliari2006}. The source can stay in the hard state to fairly high luminosities ($\sim 30 \%$ Eddington) until it changes its state. During the state transition, it first enters the \emph{hard intermediate state} (IMS), which is characterized by a hard spectral component and band-limited noise in the power spectrum. It is usually accompanied by an increasingly unstable jet \citep{FenderBelloniGallo2004}. After the hard IMS it enters the soft IMS, which is dominated by a soft spectral component and power-law noise in the power spectrum. Near this transition one often observes a bright radio flare; however, after the flare the jet seems to be quenched. After the soft IMS, the source is often found in the \emph{soft state} where the X-ray spectrum is dominated by a soft multi-colour black body component. Here the jet is typically quenched by a factor $\ga 50$ in GHz radio luminosity \citep{FenderCorbelTzioumis1999,CorbelFenderTzioumis2000}. In the soft state it is generally assumed that the source is efficiently accreting, i.e., it has a standard geometrically thin, optically thick accretion disk \citep{ShakuraSunyaev1973}. On the further way back down to low accretion rates, the source stays in the soft state until it reaches a critical accretion rate around $\sim 2\%$ of the Eddington rate \citep{Maccarone2003}, where the reverse state transition begins. This proceeds again via the soft IMS and the hard IMS to the hard state. The transition from the hard to the soft state typically occurs at a higher luminosity than the transition from the soft to the hard state. Due to this hysteresis effect there is no one-to-one correspondence of accretion rates to accretion states. For a detailed discussion of states and their exact definitions see \citet{Nowak1995,BelloniHomanCasella2005,HomanBelloni2005}, but see also \citet{McClintockRemillard2003} for slightly different definitions (mainly concerning the intermediate states). Both in X-ray binaries and AGN, the unbeamed radio luminosity can be used to estimate the accretion rate. \citet{KoerdingFenderMigliari2006} do this for the \emph{core} radio emission, while \citet{WillottRawlingsBlundell1999} present a correlation between \emph{extended} radio luminosity and tracers of the accretion rate of luminous double-lobed radio sources. Here, we use these methods to translate radio luminosity functions into accretion rate functions (ARF) of jet-emitting sources. In order to turn a \emph{radiative} bolometric luminosity function (BLF) into an ARF as well, it is necessary to know the radiative efficiency of the accretion flow. There have been several suggestions that inefficient accretion is visible both in XRBs \citep[e.g.,][]{Ichimaru1977,EsinMcClintockNarayan1997} and AGN (e.g., \citealt{ReesPhinneyBegelman1982,ChiabergeGilliMacchetto2006}), and a number of authors have attempted to link both source classes and establish a detailed correspondence (see, e.g., \citealt*{Meier2001,LivioPringleKing2003,MaccaroneGalloFender2003,MerloniHeinzdiMatteo2003,FalckeKoerdingMarkoff2004,Jester2005}; \citealt*{KoerdingFalckeCorbel2005}; \citealt{KoerdingFenderMigliari2006}). These sources are not only likely to be inefficiently accreting, but the total energy output of the sources may be dominated by the jet power \citep[e.g.,][]{FenderGalloJonker2003,KoerdingFenderMigliari2006}. To compare the ARF with the BLF we will explore the effect of inefficient accretion flows for low-luminosity objects that are included in the bolometric luminosity function of \citep{HopkinsRichardsHernquist2007}. Comparing our radio-derived ARF obtained in this way with this BLF, we can also determine a radio-loud fraction that has an intuitive physical meaning: the ratio of volume densities of radio-loud and radio-quiet sources at a given total accretion rate. As the jet power is employed in our method for estimating the accretion rates from radio data, we can also directly estimate the jets' kinetic power from the radio luminosities \citep[for an approach estimating jet powers from modeling blazar SEDs, see][]{MaraschiTavecchio2003}. This allows us to construct kinetic luminosity functions for jet emitting sources in the local universe as well as for higher redshifts. Such kinetic luminosity functions have already been constructed by \citet{HeinzMerloniSchwab2007}. However, we use a slightly different methodology and are able to include high-luminosity sources (i.e., FR-II radio galaxies). By integrating the kinetic luminosity functions over all luminosities we then estimate the total power available for kinetic and radiative feedback at a given redshift. We also consider constraints on the Eddington ratio distribution of accreting black holes which reproduce the observed luminosity functions of \emph{all} classes of AGN using only the black-hole mass function and prescriptions for the radiative and kinetic outputs of accreting black holes as function of accretion rate \citep[compare][]{VolonteriSalvaterraHaardt2006,MarulliBranchiniMoscardini2007}. | For both AGN and X-ray binaries, both the accretion rate and the jet power can be estimated from either the core radio luminosity \citep{KoerdingFenderMigliari2006} or the extended low-frequency radio luminosity \citep{WillottRawlingsBlundell1999}. Using these accretion rate estimates, we construct accretion rate functions (ARFs) of jet-emitting sources. The luminosity function (LF) of low-luminosity sources -- which likely have slow jets -- is roughly in agreement with the bolometric luminosity function (BLF) of AGN (Fig.~\ref{FigLumZ0}) as determined by \citet{HopkinsRichardsHernquist2007}. The ARF of RLQ-like jet sources is roughly 1\,dex below the BLF (Fig.~\ref{Figfrac}). This provides a measurement of the radio-loud fraction with a direct physical meaning: the ratio of volume densities of radio-loud and radio-quiet sources at a given accretion rate. However, the exact value is unfortunately strongly dependent on the normalizations used to derive accretion rates. We have developed a simple model based on the universality of accretion physics in XRBs and AGN that reproduces both our ARF as well as the BLF. At low luminosities, all sources are radiatively inefficient and the luminosity depends quadratically on the accretion rate (eq.~\ref{eqbol}). Sources are radiatively efficient only at high luminosities ($\ga 3\%$ Eddington). In our model, we assumed that the distribution of accretion rates does not depend on the black-hole mass and can be written as a simple broken power law for the Eddington-scaled accretion rate (eq.~\ref{eqmdotdist}). With this model we can reproduce the measured ARF as well as the local BLF (Fig.~\ref{fiARFQLF}). Thus, this model can solve the problem that there seem to be too few LLAGN compared to the number of weak quasars. In this simple model the break in the luminosity function is due to the inefficient accretion at low accretion rates \citep[compare][]{Jester2005}. Using the corresponding jet power measures, we have calculated kinetic luminosity functions and their cosmological evolution. Our findings support the idea that the majority of the kinetic power is created by lower-luminosity AGN with their likely slow jets. The highly luminous radio-loud quasars do not contribute significantly due to their low number density. The total power injected into the ambient medium by jets from AGN is comparable or even exceeds the effect of supernovae (the situation may and will be different for individual objects). Only at high redshifts, ``quasar-mode'' feedback mechanisms provide a comparable amount of energy for heating the interstellar/intergalactic gas. Finally, we discuss the effects of the different feedback mechanisms and find that a roughly constant fraction (5-10$\%$) of the accretion power is available for feedback, independent of the nature of the AGN. | 7 | 10 | 0710.1718 |
0710 | 0710.4417_arXiv.txt | {We present the results of a deep (1.1 Ms) observation of the Coma cluster of galaxies in the 18-30 keV band with the ISGRI imager on board the \emph{INTEGRAL} satellite. We show that the source extension in the North-East to South-West (SW) direction ($\sim 17'$) significantly exceeds the size of the point spread function of ISGRI, and that the centroid of the image of the source in the 18-30~keV band is displaced in the SW direction compared to the centroid in the 1-10~keV band. To test the nature of the SW extension we fit the data assuming different models of source morphology. The best fit is achieved with a diffuse source of elliptical shape, although an acceptable fit can be achieved assuming an additional point source SW of the cluster core. In the case of an elliptical source, the direction of extension of the source coincides with the direction toward the subcluster falling onto the Coma cluster. If the SW excess is due to the presence of a point source with a hard spectrum, we show that there is no obvious X-ray counterpart for this additional source, and that the closest X-ray source is the quasar EXO 1256+281, which is located $6.1'$ from the centroid of the excess. Finally, we show that the hard X-ray emission coincides with the 1.4 GHz radio emission, which suggests that the hard X-ray emission comes from the same population of electrons that is responsible for radio haloes through synchrotron emission.} \begin{document} | In the hierarchical scenario of structure formation, clusters of galaxies are the latest and biggest structures to form. Hence, we expect some of them to be still forming, and experiencing major merging events with smaller clusters. This is the case of the Coma cluster, that is currently merging with the NGC 4839 group. In such events, the merging of the ICM of the two clusters creates shock fronts, in which theory predicts that an important population of particles would be accelerated to high energies \citep{sarazin}. This phenomenon should then produce a reheating of the gas, and create a higher temperature plasma that would radiate more strongly in hard X-rays. Alternatively, interaction of the population of mildly relativistic electrons that produce the halos of galaxy clusters via synchrotron radiation \citep{feretti} with the Cosmic Microwave Background would then produce hard X-ray emission through inverse Compton processes, and thus add a power-law tail to the spectrum in the hard X-ray domain. Detection of this hard X-ray excess would help to learn more about the cosmic ray population detected by radio observations. Furthermore, characterization of the morphology of the hard X-ray emission would bring a possible identification of acceleration sites. Recent reports of detection of a hard X-ray excess by \emph{Beppo-SAX} \citep{fusco} and \emph{RXTE} \citep{rephaeli} in the Coma cluster seem to confirm the existence of a high energy tail of the spectrum of merging clusters, and thus prove the existence of particle acceleration sites in these clusters. However, these detections are rather weak and controversial \citep{rossetti}, and since the hard X-ray instruments on both \emph{Beppo-SAX} and \emph{RXTE} are non-imaging, contamination by very hard point sources inside the cluster could not be excluded (e.g. by the central galaxy NGC 4874, NGC 4889 or the QSO EXO 1256+281). Besides, it was not possible to have any information on the morphology of the hard X-ray emission. | Thanks to the imaging capabilities of ISGRI, we were able for the first time to resolve spatially the Coma cluster in the hard X-ray domain. We showed that the hard X-ray emission is displaced compared to the purely thermal emission in soft X-rays. Investigating this displacement, we found that the contribution of point sources to the observed spectrum is likely negligible. Our analysis shows that a hotter plasma can explain this excess. However, the correlation between the radio and hard X-ray morphology strongly suggests that the SW excess is due to IC scattering of mildly relativistic electrons. Together with the spectral results obtained by \emph{Beppo-SAX} \citep{fusco2}, it is now becoming clear that the presence of an additional X-ray spectral component is required by the data. Moreover, the strong correlation between radio and hard X-ray morphology clarifies the nature of this excess, and seems to confirm the existence of a particle acceleration site in the Coma cluster. | 7 | 10 | 0710.4417 |
0710 | 0710.1697_arXiv.txt | {A large fraction of otherwise similar asymptotic giant branch stars (AGB) do not show OH maser emission. As shown recently, a restricted lifetime may give a natural explanation as to why only part of any sample emits maser emission at a given epoch.} % {We wish to probe the lifetime of 1612 MHz OH masers in circumstellar shells of AGB stars.} {We reobserved a sample of OH/IR stars discovered more than 28 years ago to determine the number of stars that may have since lost their masers.} {We redetected all 114 OH masers. The minimum lifetime inferred is 2800 years (1$\sigma$). This maser lifetime applies to AGB stars with strong mass loss leading to very red infrared colors. The velocities and mean flux density levels have not changed since their discovery. As the minimum lifetime is of the same order as the wind crossing time, strong variations in the mass-loss process affecting the excitation conditions on timescales of $\approx$3000 years or less are unlikely.} {} | Molecular OH maser emission at 1612 MHz is frequently observed from stars approaching the tip of their evolutionary track on the asymptotic giant branch (AGB). These stars are in general long-period variables with pulsation periods $>$1~year. The pulsation is accompanied by heavy mass loss, which forms a circumstellar envelope (CSE) of gas and dust, and if the mass-loss rate surpasses \mdot\ $\approx$10$^{-6}$ \Myr\ the dust envelope becomes opaque for visible light. Classically these `invisible' stars were called OH/IR stars, but nowadays this term is often used in general reference to OH emitting AGB stars selected in the infrared (see review by Habing \cite{Habing96}). Interestingly, a large fraction ($\approx$40\%) of any sample of IRAS sources, selected to match the infrared colors of known OH/IR stars, does not exhibit a detectable OH 1612 MHz maser (Lewis \cite{Lewis92}). These infrared sources were baptized by Lewis as `OH/IR star color mimics'. Some of them do exhibit mainline OH (at 1665 or 1667 MHz) and/or 22 GHz \water\ masers (Lewis \& Engels \cite{Lewis95}), corroborating their cousinhood with OH/IR stars. OH/IR stars are therefore only part of the population of evolved and obscured AGB stars with oxygen-rich chemistry. The reasons for the absence of 1612 MHz OH masers in a large fraction of oxygen-rich AGB stars are not known. The maser photons are emitted by a transition between hyperfine levels of the groundstate of OH, which are inverted with the help of pump photons at $\lambda$=35 and 53\micron, emitted from dust in the CSEs. A requirement for excitation is sufficient OH column density, which might be low in mimics due to the destruction of OH by the interstellar UV field. This effect cannot account for mimics in general, as mimics are also present at higher galactic latitudes, where the influence of the UV field is low. In addition, the presence of a hot white dwarf companion leading to photodissociation of OH is not generally able to suppress the OH maser (Howe \& Rawlings \cite{Howe94}). A further requirement is velocity coherence over large distances to allow amplification. Turbulence may disrupt this coherence in the case of mimics. An alternative explanation is the assumption that the OH maser is only present temporarily on the AGB and therefore these stars may change between an OH/IR status and that of a mimic. This explanation has been triggered by recent observations showing the fading of the OH maser in \object{IRAS~18455+0448} over a time range of a decade (Lewis et al. \cite{Lewis01}) and the absence of masers in four stars during a revisit of 328 OH/IR stars 12 years after their first detection (Lewis \cite{Lewis02}). The high rate of `dead' OH/IR stars among a sub-sample of 112 OH/IR stars with relatively blue colors led Lewis to conclude that the mean 1612 MHz emission life is in the range 100--400 years. To test this lifetime we reobserved another sample of N$>$ 100 OH/IR stars, which was drawn from the first surveys for OH maser emission prior to 1980. With a difference of almost 30 years between the two observations several masers were expected to have disappeared. | We find that the lifetime of OH masers in classical OH/IR stars is $>2800$ years, in contrast to the lifetime implied by the observed disappearance of several masers in AGB stars with bluer envelopes. The lower limit of the lifetime is of the order of the wind crossing time in the CSEs, implying that, in general, no drastic variations in the structure of the CSEs happen on such timescales, or shorter ones. If non-variable OH/IR stars are in transition to the post-AGB stage, a sudden decline of the mass-loss rates at the end of AGB evolution is ruled out. The previous observed disappearance of OH masers in stars with bluer envelopes is probably not associated with major changes in the mass-loss process. The susceptibility of masers to smaller variances in their environment may lead to OH masers as transient phenomena on the AGB. This gives a natural explanation for the detection of OH masers in only a part of any AGB star sample with otherwise similar properties. | 7 | 10 | 0710.1697 |
0710 | 0710.1142_arXiv.txt | {} % {We present new millimetre 43 GHz observations of a sample of radio-bright Planetary Nebulae. Such observations were carried out to have a good determination of the high-frequency radio spectra of the sample in order to evaluate, together with far-IR measurements (IRAS), the fluxes emitted by the selected source in the millimetre and sub-millimetre band. This spectral range, even very important to constraint the physics of circumstellar environment, is still far to be completely exploited.} {To estimate the millimetre and sub-millimetre fluxes, we extrapolated and summed together the ionized gas (free-free radio emission) and dust (thermal emission) contributions at this frequency range. By comparison of the derived flux densities to the foreseen sensitivity we investigate the possible detection of such source for all the channels of the forthcoming ESA's PLANCK mission.} {We conclude that almost 80\% of our sample will be detected by PLANCK, with the higher detection rate in the higher frequency channels, where there is a good combination of brighter intrinsic flux from the sources and reduced extended Galactic foregrounds contamination despite a worst instrumental sensitivity. From the new 43 GHz, combined with single-dish 5 GHz observations from the literature, we derive radio spectral indexes, which are consistent with optically thin free-free nebula. This result indicates that the high frequency radio spectrum of our sample sources is dominated by thermal free-free and other emission, if present, are negligible. } {} % | The PLANCK ESA mission will provide us with nearly full sky maps over a wide range of frequencies, from 30 to 900 GHz. Therefore, even if designed for cosmological studies, the mission will have profund impact on fundamental Physics and Galactic and extragalactic astrophysics. Planck will be sensitive to the millimetre emission from dusty envelopes of stars and we expect the dusty circumstellar envelopes, that characterize the latest stages of stellar evolution, to be source of relevant foreground contamination. When a low or intermediate mass star is approaching the end of its evolution, it goes through a period of heavy mass-loss known as Asymptotic Giant Branch (AGB) phase. The ejected envelopes are partially condensed in dust grains and completely obscure the central star. Immediately after the AGB evolutionary phase, the mass-loss stops and the stars may become optically visible as the dusty shells disperse (Proto Planetary Nebula, PPN phase). Eventually, once it reaches a temperature of 2 -- 3 $\times 10 ^4$ K, the central star starts to ionize the AGB shell and a Planetary Nebula will form. Dusty envelopes re-radiate the absorbed stellar light showing a clear signature in the far-infrared spectrum, i.e. an IR excess with peculiar IRAS colours. In addition to that, PNe show a radio continuum due to free-free emission from the fraction of the CSE ionized by the central star. \begin{table*} \caption{The selected sample} \label{tab_posizioni} \begin{center} \vspace{0.3cm} \begin{tabular}{cccc|cccc} \hline \hline IAU Name & Other Name & R.A.\ (J2000) &Dec. \ (J2000) &IAU Name & Other Name & R.A.\ (J2000) & Dec. \ (J2000) \\ PN G&& [\ h\ m\ s] &[$^{\circ}\ ^{\prime}\ ^{\prime\prime}$] &PN G&&[\ h\ m\ s] &[$^{\circ}\ ^{\prime}\ ^{\prime\prime}$] \\ \hline $000.3+12.2$ &IC 4634 & 17 01 33.6 & $-21$ 49 33.1 &$093.4+05.4$ &NGC 7008 & 21 00 32.7 & $+54$ 32 39.4 \\ $002.4+05.8$ &NGC 6369 & 17 29 20.5 & $-23$ 45 35.0 &$093.5+01.4$ &PN M 1-78 & 21 20 44.8 & $+51$ 53 27.5 \\ $003.1+02.9$ &PN Hb 4 & 17 41 52.8 & $-24$ 42 09.3 &$096.4+29.9$ &NGC 6543 & 17 58 33.4 & $+66$ 37 58.8 \\ $006.7-02.2$ &PN M 1-41 & 18 09 30.6 & $-24$ 12 28.7 &$097.5+03.1$ &PN A66 77 & 21 32 10.2 & $+55$ 52 43.2 \\ $007.2+01.8$ &PN HB 6 & 17 55 07.0 & $-21$ 44 41.0 &$106.5-17.6$ &NGC 7662 & 23 25 53.9 & $+42$ 32 04.7 \\ $008.0+03.9$ &NGC 6445 & 17 49 15.0 & $-20$ 00 33.7 &$107.8+02.3$ &NGC 7354 & 22 40 19.9 & $+61$ 17 08.0 \\ $008.3-01.1$ &PN M 1-40 & 18 08 26.0 & $-22$ 16 53.4 &$120.0+09.8$ &NGC 40 & 00 13 01.0 & $+72$ 31 19.6 \\ $009.4-05.0$ &NGC 6629 & 18 25 42.5 & $-23$ 12 11.3 &$130.9-10.5$ &NGC 650-51 & 01 42 19.7 & $+51$ 34 31.7 \\ $009.6+14.8$ &NGC 6309 & 17 14 04.3 & $-12$ 54 37.2 &$138.8+02.8$ &IC 289 & 03 10 19.3 & $+61$ 19 00.4 \\ $010.1+00.7$ &NGC 6537 & 18 05 13.1 & $-19$ 50 34.4 &$144.5+06.5$ &NGC 1501 & 04 06 59.3 & $+60$ 55 14.7 \\ $010.8-01.8$ &NGC 6578 & 18 16 16.5 & $-20$ 27 03.4 &$165.5-15.2$ &NGC 1514 & 04 09 16.9 & $+30$ 46 32.0 \\ $011.7-00.6$ &NGC 6567 & 18 13 45.2 & $-19$ 04 35.6 &$166.1+10.4$ &IC 2149 & 05 56 23.9 & $+46$ 06 17.4 \\ $020.9-01.1$ &PN M 1-51 & 18 33 29.0 & $-11$ 07 26.3 &$173.7+02.7$ &PP 40 & 05 40 52.7 & $+35$ 42 18.6 \\ $025.8-17.9$ &NGC 6818 & 19 43 57.8 & $-14$ 09 11.8 &$194.2+02.5$ &J 900 & 06 25 57.3 & $+17$ 47 27.6 \\ $027.7+00.7$ &PN M 2-45 & 18 39 21.9 & $-04$ 19 52.6 &$197.8+17.3$ &NGC 2392 & 07 29 10.8 & $+20$ 54 41.6 \\ $033.8-02.6$ &NGC 6741 & 19 02 37.0 & $-00$ 26 57.2 &$206.4-40.5$ &NGC 1535 & 04 14 15.8 & $-12$ 44 22.3 \\ $034.6+11.8$ &NGC 6572 & 18 12 06.3 & $+06$ 51 12.4 &$215.2-24.2$ &IC 418 & 05 27 28.2 & $-12$ 41 50.2 \\ $035.1-00.7$ &PN Ap 2-1 & 18 58 10.5 & $+01$ 36 57.5 &$221.3-12.3$ &IC 2165 & 06 21 42.8 & $-12$ 59 13.9 \\ $037.7-34.5$ &NGC 7009 & 21 04 10.8 & $-11$ 21 48.5 &$234.8+02.4$ &NGC 2440 & 07 41 55.4 & $-18$ 12 30.5 \\ $039.8+02.1$ &PN K 3-17 & 18 56 18.2 & $+07$ 07 26.2 &$254.6+00.2$ &NGC 2579 & 08 20 54.1 & $-36$ 13 00.0 \\ $041.8-02.9$ &NGC 6781 & 19 18 28.1 & $+06$ 32 20.0 &$258.1-00.3$ &Hen 2-9 & 08 28 28.0 & $-39$ 23 39.4 \\ $043.1+37.7$ &NGC 6210 & 16 44 29.5 & $+23$ 47 59.9 &$259.1+00.9$ &Hen 2-11 & 08 37 08.1 & $-39$ 25 04.9 \\ $045.7-04.5$ &NGC 6804 & 19 31 35.1 & $+09$ 13 30.2 &$261.0+32.0$ &NGC 3242 & 10 24 46.1 & $-18$ 38 32.3 \\ $050.1+03.3$ &PN M 1-67 & 19 11 31.1 & $+16$ 51 32.0 &$294.1+43.6$ &NGC 4361 & 12 24 30.8 & $-18$ 47 04.0 \\ $054.1-12.1$ &NGC 6891 & 20 15 08.9 & $+12$ 42 15.4 &$342.1+10.8$ &NGC 6072 & 16 12 58.4 & $-36$ 13 46.6 \\ $063.1+13.9$ &NGC 6720 & 18 53 35.1 & $+33$ 01 45.1 &$349.5+01.0$ &NGC 6302 & 17 13 44.5 & $-37$ 06 11.6 \\ $064.7+05.0$ &BD+30 3639 & 19 34 45.2 & $+30$ 30 59.2 &$352.6+00.1$ &PN H 1-12 & 17 26 24.3 & $-35$ 01 41.8 \\ $082.1+07.0$ &NGC 6884 & 20 10 23.7 & $+46$ 27 40.0 &$352.8-00.2$ &PN H 1-13 & 17 28 27.7 & $-35$ 07 30.4 \\ $083.5+12.7$ &NGC 6826 & 19 44 48.2 & $+50$ 31 31.3 &$358.5+02.6$ &PN HDW 8 & 17 31 47.3 & $-28$ 42 03.5 \\ $086.5-08.8$ &PN Hu 1-2 & 21 33 08.2 & $+39$ 38 08.3 &$358.5+05.4$ &PN M 3-39 & 17 21 11.5 & $-27$ 11 37.0 \\ $089.0+00.3$ &NGC 7026 & 21 06 18.7 & $+47$ 51 07.5 &$359.3-00.9$ &Pn HB 5 & 17 47 56.3 & $-29$ 59 40.6 \\ \hline \end{tabular} \end{center} \end{table*} From a feasibility study, Umana et al. (\cite{Umana06}) concluded that a sizable ($\approx 300$) sample of AGB and post-AGB stars would be detected during the mission and derived estimates for the expected flux densities at various Planck channels. However, the simulations carried out by Umana et al. (\cite{Umana06}), on a sub-sample of PNe rely only on NVSS fluxes, obtained at 1.4 GHz, extrapolated to the Planck frequencies. This leads to underestimate the contribution due to free-free as, at this frequency, PNe are often optically thick (Siodmiank \& Tylenda \cite{st01}, Luo et al. \cite{luo_etal05}). PNe are among the brightest Galactic radio sources. Some of them could also reach a flux density of Jy level. More than 800 have been detected at least at one frequency and for 200 morphological and spectral information (between 1.4 and 22 GHz) were obtained. Most of the higher frequency (22 GHz) data were obtained with interferometers (i.e. VLA) while very little single-dish, high frequency measurements are available. In this paper we present new 43 GHz, single-dish observations by using the 32 m INAF-IRA Radiotelescope at Noto of a sample of PNe, which are potential foregrounds for PLANCK. The main goal of this project is to obtain reliable estimates of flux density expected at PLANCK channels, by building and modeling their spectral energy distribution (SED). This in turn will contribute to the compilation of the PLANCK pre-launch catalogue. We stress here that 43 GHz is one of the observing channels of the forthcoming PLANCK mission. Therefore, at this frequency band, we would obtain a direct measurement of the expected flux and not an extrapolation. As added value, the present work provides the first sizeable dataset of 43 GHz measurements of PNe, that constitute strong constraints to the observed SEDs in the very important spectral region where free-free emission and thermal dust emission may overlap. While interferometric high frequency observations provide us with detailed morphological information they quite often fail to entirely recover the extended emission. This eventually leads to underestimate the total radio flux density. This problem is overcame by single-dish observations. Assesting the SEDs in the radio-millimetre spectral range is a crucial step for the study of the physics of dusty envelopes around PN. A correct evaluation of the free-free contribution, up to millimetre range, when combined with information provide by far-IR observations, would allow to determine the presence of an excess due to the presence of a cold dust component/s (Gee et al. \cite{gee_etal84}; Hoare et al. \cite{hoare_etal92}; Kemper et al. \cite{kemper_etal02}) or of alternative emission mechanisms (Casassus et al. \cite{cas_etal07}). Since PNe and their progenitors are believed to be among the major sources of recycled interstellar matter, determining the properties of the dust ejected in the ISM is very important to the study of the Galaxy evolution in general. \begin{table*} \begin{center} \caption{Results} \label{tab_misure} \begin{tabular}{crrclr|crrclr} \hline \hline IAU Name & $S_\mathrm{43 ~GHz}$ & $\sigma_\mathrm{43 ~GHz}$ &$\theta_\mathrm{1.4 ~GHz}$ & Ref.$^{\ast}$ &$S_\mathrm{c~43 ~GHz}$ &IAU Name &$S_\mathrm{43 ~GHz}$ & $\sigma_\mathrm{43 ~GHz}$ &$\theta_\mathrm{1.4 ~GHz}$ & Ref.$^{\ast}$ &$S_\mathrm{c~43 ~GHz}$ \\ PN G & [mJy] & [mJy] & [arcsec] & & [mJy] &PN G & [mJy] & [mJy] & [arcsec] & & [mJy] \\ \hline $000.3+12.2$ & $<$ 180 & 60 & 10.0 & C H & &$093.4+05.4$ & $<$ 240 & 80 & 49.2 & A E & \\ $002.4+05.8$ & 1330 & 110 & 20.7 & C H &1530 &$093.5+01.4$ & 560 & 35 & 14.6 & A E &600 \\ $003.1+02.9$ & 100 & 20 & 11.6 & C H &105 &$096.4+29.9$ & 400 & 70 & 11.9 & A E &420 \\ $006.7-02.2$ & 300 & 90 & 21.0 & &345 &$097.5+03.1$ & 230 & 25 & 33.4 & A E &320 \\ $007.2+01.8$ & 310 & 50 & 10.6 & C H &320 &$106.5-17.6$ & 610 & 40 & 12.0 & A E &640 \\ $008.0+03.9$ & 180 & 20 & 34.6 & C G &250 &$107.8+02.3$ & 280 & 25 & 16.4 & A E &305 \\ $008.3-01.1$ & 90 & 20 & 9.9 & C H &95 &$120.0+09.8$ & 420 & 70 & 27.6 & A E &530 \\ $009.4-05.0$ & 130 & 30 & 12.8 & C H &135 &$130.9-10.5$ & 140 & 40 & 59.3 & A E &310 \\ $009.6+14.8$ & 270 & 60 & 14.3 & C H &290 &$138.8+02.8$ & $<$ 150 & 50 & 23.1 & A E & \\ $010.1+00.7$ & 400 & 60 & 11.8 & C F H &420 &$144.5+06.5$ & 215 & 30 & 34.9 & A E &305 \\ $010.8-01.8$ & 130 & 20 & 11.9 & C H &135 &$165.5-15.2$ & 220 & 30 & 94.4 & A E &890 \\ $011.7-00.6$ & $<$ 150 & 50 & 8.7 & C H & &$166.1+10.4$ & 100 & 30 & 9.1 & A E &105 \\ $020.9-01.1$ & 420 & 90 & 13.5 & C &445 &$173.7+02.7$ & 260 & 20 & 11.5 & E &270 \\ $025.8-17.9$ & 270 & 30 & 14.7 & C H &290 &$194.2+02.5$ & 90 & 20 & 0.0 & A E H &90 \\ $027.7+00.7$ & 110 & 20 & 0.0 & B G &110 &$197.8+17.3$ & 240 & 50 & 22.2 & A E H &280 \\ $033.8-02.6$ & 290 & 50 & 14.7 & E H &310 &$206.4-40.5$ & 160 & 20 & 20.6 & C H &180 \\ $034.6+11.8$ & 1220 & 100 & 11.9 & A E H &1280 &$215.2-24.2$ & 1100 & 100 & 0.0 & C H &1100 \\ $035.1-00.7$ & 170 & 20 & 13.0 & A E &180 &$221.3-12.3$ & 350 & 50 & 11.0 & C H &365 \\ $037.7-34.5$ & 375 & 40 & 14.9 & C H &400 &$234.8+02.4$ & 350 & 20 & 16.5 & C H &380 \\ $039.8+02.1$ & 240 & 20 & 0.0 & A E &240 &$254.6+00.2$ & 1770 & 290 & 41.2 & F &2800 \\ $041.8-02.9$ & 230 & 20 & 78.4 & A E H &715 &$258.1-00.3$ & 180 & 40 & 4.5 & D H &181 \\ $043.1+37.7$ & 230 & 30 & 10.8 & A E H &240 &$259.1+00.9$ & 270 & 60 & 42.2 & D &435 \\ $045.7-04.5$ & 140 & 20 & 25.4 & A B E H &170 &$261.0+32.0$ & 290 & 30 & 19.5 & C H &330 \\ $050.1+03.3$ & 140 & 30 & 41.4 & A E H &220 &$294.1+43.6$ & 130 & 20 & 47.3 & C H &230 \\ $054.1-12.1$ & 130 & 40 & 8.9 & A E H &135 &$342.1+10.8$ & $<$ 210 & 70 & 32.9 & H & \\ $063.1+13.9$ & 255 & 75 & 48.2 & A E &460 &$349.5+01.0$ & 2150 & 220 & 14.8 & D F H &2310 \\ $064.7+05.0$ & 565 & 20 & 10.5 & A E &585 &$352.6+00.1$ & 815 & 125 & 10.3 & F &845 \\ $082.1+07.0$ & 250 & 50 & 11.0 & A E &260 &$352.8-00.2$ & 420 & 50 & 14.8 & F &450 \\ $083.5+12.7$ & 320 & 40 & 15.9 & A E &350 &$358.5+02.6$ & $<$ 120 & 40 & 21.1 & C G & \\ $086.5-08.8$ & $<$ 120 & 40 & 10.4 & A E & &$358.5+05.4$ & 380 & 40 & 0.0 & C G &380 \\ $089.0+00.3$ & 220 & 30 & 11.9 & A E &230 &$359.3-00.9$ & 290 & 70 & 0.0 & F H &290 \\ \hline \end{tabular} \begin{list}{}{} \item[$^{\ast}$] References for 5 GHz measurements: A) Gregory et al. (\cite{g_etal96}); B) Griffith et al. (\cite{g_etal95}); C) Griffith et al. (\cite{g_etal94}); D) Wright et al. (\cite{w_etal94}); E) Becker et al. (\cite{b_etal91}); F) Haynes et al. (\cite{h_etal79}); G) Milne (\cite{m79}); H) Milne (\cite{ma75}) \end{list} \end{center} \end{table*} | We have presented new 7 mm (43 GHz) observations of a sample of radio-bright PNe, carried out with the INAF-IRA Noto Radiotelescope. Such observations have been used to derive the high-frequency free-free contribution, due to the ionized fraction of the circumstellar envelope. So far, the majority of high frequency radio observations of PNe have been conducted with interferometers, that reveal details of source radio morphology but could, in principle, resolve out some of the extended emission. Our single dish measurements provide an extended database of millimetre observations of PNe, to be used when is necessary to know the overall emission, such as when building a SED. We used our measurements, that trace the free-free contribution, together with IRAS measurements, which trace the thermal dust emission, to build up the observed SED, from radio to far-IR and by extrapolation, to estimate the expected fluxes in the spectral range between radio and sub-mm, where observational data are missing. When comparing the expected millimetre-sub-millimetre fluxes with the total foreseen sensitivity of the forthcoming ESA mission PLANCK we estimate that a consistent number of our targets will be easily detected by PLANCK, mostly at higher frequency channels. Therefore, even if designed for cosmological study, PLANCK could also contribute to the PNe science. Results from such kinds of observations, once modelled with appropriate code (i.e. CLOUDY, DUSTY), would provide important constraints on the chemical composition and structure of the CSEs. It would point out the presence of multi-shells, related to multiple mass-loss event suffered from the central object during its previous evolution (AGB), or a contribution of different emission processes, besides free-free and thermal from dust, as recently claimed by Cassassus et al.~(\cite{cas_etal07}). Moreover, PLANCK results can be considered as pathfinder for other future instrumentations, with the same frequency coverage, such as ALMA, as they will help in planning more focused experiments aimed to investigate the morphological details of the sources. | 7 | 10 | 0710.1142 |
0710 | 0710.4883_arXiv.txt | Knowledge of the stellar parameters for the parent stars of transiting exoplanets is pre-requisite for establishing the planet properties themselves, and often relies on stellar evolution models. GJ~436, which is orbited by a transiting Neptune-mass object, presents a difficult case because it is an M dwarf. Stellar models in this mass regime are not as reliable as for higher mass stars, and tend to underestimate the radius. Here we use constraints from published transit light curve solutions for GJ~436 along with other spectroscopic quantities to show how the models can still be used to infer the mass and radius accurately, and at the same time allow the radius discrepancy to be estimated. Similar systems should be found during the upcoming \emph{Kepler\/} mission, and could provide in this way valuable constraints to stellar evolution models in the lower main sequence. The stellar mass and radius of GJ~436 are $M_{\star} = 0.452_{-0.012}^{+0.014}$~M$_{\sun}$ and $R_{\star} = 0.464_{-0.011}^{+0.009}$~R$_{\sun}$, and the radius is 10\% larger than predicted by the standard models, in agreement with previous results from well studied double-lined eclipsing binaries. We obtain an improved planet mass and radius of $M_p = 23.17 \pm 0.79$~M$_{\earth}$ and $R_p = 4.22_{-0.10}^{+0.09}$~R$_{\earth}$, a density of $\rho_p = 1.69_{-0.12}^{+0.14}$ g~cm$^{-3}$, and an orbital semimajor axis of $a = 0.02872 \pm 0.00027$ AU. | \label{sec:introduction} The applications of stellar evolution theory to astrophysics are so widespread, and its validity so often taken for granted, that it is easy to forget that it took decades to develop, and significant effort to validate by comparison with careful measurements, a process that still continues. It is usually only when those theoretical predictions fail that the classical discipline of stellar evolution ``makes the headlines'', and even then it draws the attention of relatively few. One such instance has occurred for low mass main-sequence stars. Over the last 10 years or so it has become clear that our understanding of the structure and evolution of these objects is still incomplete. Discrepancies between theory and observation in the radii of stars under 1~M$_{\sun}$, first mentioned by \cite{Hoxie:73}, \cite{Lacy:77}, and others, are now well documented for several low-mass eclipsing binaries \citep[see, e.g.,][]{Popper:97, Clausen:99, Torres:02, Ribas:03, Lopez-Morales:05}. Differences in the effective temperatures have been observed as well. The direction of these disagreements is such that model radii are underestimated by roughly 10\%, while effective temperatures are overestimated. In recent years stellar evolution has had important applications in the field of transiting extrasolar planets. This is because the planetary parameters of interest (mass $M_p$, radius $R_p$) depend rather directly on those of the star ($M_{\star}$, $R_{\star}$), and in most cases models provide the only means of determining the latter. The subject of this paper is GJ~436, a late-type star found by \cite{Butler:04} to be orbited by a Neptune-mass planet with a period of 2.644 days. This object was later discovered by \cite{Gillon:07a} to undergo transits, enabling its size to be determined ($\sim$4~R$_{\earth}$). As the only M dwarf among the 22 currently known transiting planet host stars, GJ~436 (M2.5V) presents a special challenge for establishing the stellar parameters, because of the disagreements noted above. Little mention seems to have been made of this, and for the most part past studies have relied instead on empirical mass-luminosity ($M\!-\!L$) relations to set the mass of GJ~436. Radius estimates have often rested on the assumption of numerical equality between $M_{\star}$ and $R_{\star}$ for M stars. Despite being the closest transiting planet system (only 10 pc away), it is rather surprising that the mass of the star is only known to about 10\% \citep{Maness:07, Gillon:07b}. This is currently limiting the precision of the planetary mass, and some of that uncertainty translates also to the radius. These two properties are critical for studying the structure of the object. Given the importance of GJ~436 as the parent star of the only Neptune-mass transiting exoplanet found so far, and hence the closest analog to our Earth with a mass and radius determination, one of the motivations of this paper is to improve the precision of the stellar and planetary parameters by making use of additional observational constraints not used before. Specifically, we incorporate the information on the stellar \emph{density} directly available from the transit light curve \citep{Sozzetti:07}, which provides a strong handle on the size of the star. The nature of the discrepancies between evolutionary models and observations for low-mass stars has been examined recently from both the observational and theoretical points of view by \cite{Lopez-Morales:07} and \cite{Chabrier:07}. Further progress depends on gathering more evidence to supplement the few available highly accurate mass and radius measurements based on double-lined eclipsing binaries. Thus, a second motivation for this work, despite the fact that GJ~436 is not a double-lined eclipsing binary, is to present a way of using all observational constraints simultaneously to show that the star presents the same radius anomaly found for the other systems, or more generally, to test the models. Because similar constraints may become available in the future for other M dwarfs given the keen interest in finding smaller and smaller transiting planets, we anticipate that the indirect technique described here may yield valuable information on this problem and eventually help improve our understanding of low-mass stars. | With the host star properties known, the planet parameters we infer are given in the bottom section of Table~\ref{tab:properties}. The required orbital period and velocity semi-amplitude $K_{\star}$ are adopted from \cite{Maness:07}, along with the eccentricity from \cite{Demory:07}, who obtain a similar value for $K_{\star}$. The improved precision of these derived parameters is a reflection of the better stellar parameters. The slightly larger planet radius than in previous studies confirms with even greater statistical significance the conclusions of earlier authors regarding the presence of a hydrogen/helium envelope \citep{Gillon:07a, Deming:07, Gillon:07b}, and agrees very well with the models by \cite{Fortney:07} for a 10\% fraction of those elements. In this paper we have shown that GJ~436 provides a valuable test of stellar evolution theory near the bottom of the main sequence, made possible by the fact that it has a transiting planet. Traditional studies in the area of low-mass stars have made the comparison with models by measuring the mass and radius directly for double-lined eclipsing systems containing M dwarfs. In a few other cases angular diameters have been measured interferometrically for single stars, and the mass has been inferred from empirical $M\!-\!L$ relations \citep[e.g.,][]{Lane:01, Segransan:03}. More recently, a variety of constraints and assumptions have been used to infer the mass and radius of the late-type secondaries in F+M systems observed as part of transiting planet surveys \citep[e.g.,][]{Bouchy:05, Pont:05, Beatty:07}. Though perhaps not as compelling as having actual model-independent mass and radius measurements, the approach in the present work is able to make use of available information for GJ~436 and compare the models directly with the observational constraints \emph{without} requiring a direct measurement of the mass and radius. The discrepancy in $R_{\star}$ is derived by parameterizing it in terms of a single adjustment factor to the model radii ($\beta$), assuming the luminosity from theory is accurate, as other observations seem to indicate. NASA's upcoming \emph{Kepler\/} mission, currently slated to launch in early 2009, will emphasize the search for transiting Earth-size planets. These should be easier to detect around late-type stars. Therefore, we anticipate that many systems similar to GJ~436 could be found and become a significant source of information on radii for low-mass stars, since they will have all the observational constraints needed (including trigonometric parallaxes) to test models of stellar evolution in the way we have done here. | 7 | 10 | 0710.4883 |
0710 | 0710.0372_arXiv.txt | We present a sample of \total\ galaxy nuclei from 12 nearby ($z<4500$\kms) Hickson Compact Groups (HCGs) with a complete suite of 1--24\micron\ 2MASS+\spitzer\ nuclear photometry. For all objects in the sample, blue emission from stellar photospheres dominates in the near-infrared through the 3.6\micron\ IRAC band. Twenty-five of \total\ (54\%) galaxy nuclei show red, mid-infrared continua characteristic of hot dust powered by ongoing star formation and/or accretion onto a central black hole. We introduce \alphairac, the spectral index of a power-law fit to the 4.5--8.0\micron\ IRAC data, and demonstrate that it cleanly separates the mid-infrared active and non-active HCG nuclei. This parameter is more powerful for identifying low to moderate-luminosity mid-infrared activity than other measures which include data at rest-frame $\lambda<3.6$\micron\ that may be dominated by stellar photospheric emission. While the HCG galaxies clearly have a bimodal distribution in this parameter space, a comparison sample from the \spitzer\ Nearby Galaxy Survey (SINGS) matched in $J$-band total galaxy luminosity is continuously distributed. A second diagnostic, the fraction of 24\micron\ emission in excess of that expected from quiescent galaxies, \fracd, reveals an additional three nuclei to be active at 24\micron. Comparing these two mid-infrared diagnostics of nuclear activity to optical spectroscopic identifications from the literature reveals some discrepancies, and we discuss the challenges of distinguishing the source of ionizing radiation in these and other lower luminosity systems. We find a significant correlation between the fraction of mid-infrared active galaxies and the total \HI\ mass in a group, and investigate possible interpretations of these results in light of galaxy evolution in the highly interactive system of a compact group environment. | \label{sec:intro} The first galaxies and their environments differed substantially from those locally, often involving multiple interactions as seen in the \hst\ Ultra Deep Field \citep[e.g.,][]{malhotra05}. Compared to all other nearby environments, present-day compact galaxy groups most closely reproduce the interaction environment of the early Universe ($z\sim4$) when galaxies assembled through hierarchical formation \citep[e.g.,][]{baron}, and galaxy groups combined to form proto-clusters (in dense regions; e.g., \altcite{rudick+06}) or massive ellipticals (in the field; \altcite{white}). Because of their high space densities (with comparable surface densities to the centers of rich galaxy clusters; e.g., \altcite{rubin+91}) and low velocity dispersions ($\sigma_{\rm v}\sim10^2$~\kms), compact groups of galaxies are ideal environments for studying the mechanisms of interaction-induced star formation and nuclear activity. From optical spectroscopic surveys, Hickson Compact Groups (HCGs) are known to host a population of galaxies with emission-line nuclear spectra characteristic of star-formation and/or active galactic nuclei (AGNs). Based on the optical spectroscopic survey of \citet{coziol98a,coziol98b}, the AGN fraction in HCGs is found to be $\sim40\%$, perhaps consistent with the $43\%$ nuclear activity level found for nearby $M_V<-19$ galaxies (with greater detection sensitivity; \altcite{HoEtal97a}) and significantly higher than the $\sim1\%$ AGN fraction identified optically in cluster galaxies (with $M_{V}<-21$; \altcite{dressler85}). Further, HCG AGNs (including low-luminosity AGNs, hereafter LLAGNs; HCGs host no known Seyfert~1-luminosity AGNs) are preferentially found in optically luminous, early-type galaxies with little or no ongoing star formation in the cores of evolved groups. Similarly, many galaxies in clusters host LLAGNs in the local Universe \citep[e.g.,][]{martini07}, in particular brightest cluster galaxies \citep{best07}. Typically, the most active star-forming galaxies are late-type spirals in the outskirts of groups. \citet{coziol98b} interpreted relative distributions of star-forming and AGN-hosting galaxies as indicating a clear evolutionary scenario whereby group cores are mature collapsed systems in which the high galaxy densities led to increased gravitational interactions and hence more rapid exhaustion of gas reservoirs in the past through star formation. (The radiatively weak AGNs in the evolved group cores require only a small amount of gas for fueling.) A similar study by \citet{shimada00} found that the fraction of emission-line HCG galaxies was comparable to the field, and conversely concluded that the HCG environment does {\it not} trigger either star formation or AGN activity. The initial expectation that the interactions evident in the compact group environment would naturally lead to markedly enhanced levels of star formation compared with the field has not been satisfied, and a coherent understanding of the history of gas and cold dust in compact groups has proven elusive. An analysis of the CO content in HCG galaxies found them to be similar to those in the field, in loose groups, and in other environments; a notable exception is the $\sim20\%$ of HCG spirals that are CO deficient. This suggests less, not more, star formation in HCGs compared to other environments \citep{verdes98}. At the same time, the detection of a few HCG elliptical and S0 galaxies in both CO and the far-infrared (unlike typical galaxies of these types) suggests that tidal interactions are influencing galaxy evolution to some extent. While the far-infrared power is similar to comparison samples, the ratios of 25 to 100\micron\ \iras\ fluxes implies a greater number of intense, nuclear starbursts in HCG galaxies \citep{verdes98}. Clearly, a robust and consistent understanding of the impact of the compact group environment on galaxy properties has not yet emerged. One possible difficulty to date is the predominant use of optical emission-line studies to identify activity --- both star-forming and accretion-dominated. Ground-based studies of bright galaxies can easily obscure low-contrast emission lines through their dilution by a strong stellar contribution \citep[e.g.,][]{HoEtal97c}; intrinsic absorption can also mask spectroscopic signatures. The clear discrepancy between the $\sim1\%$ AGN fraction in galaxy clusters from optical spectroscopic surveys \citep[e.g.,][]{dressler85} compared to the larger fraction ($\sim5\%$ for luminous galaxies) revealed by X-ray observations \citep[e.g.,][]{martini06} is illustrative of this problem. As highlighted by \citet{martini07}, mismatched sample selection and detection techniques can create apparent (and false) discrepancies in AGN fractions between environments. In this paper, we take an alternate approach, focusing on the mid-infrared spectral energy distributions (SEDs) of individual galaxy nuclei to clearly identify the thermal, hot dust continua that signify neutral gas heated by ionizing photons from either young stars or an AGN. The clear discrepancy between a blue, quiescent galaxy SED where the mid-infrared is dominated by the Raleigh-Jeans tail of stellar photospheres and the red, mid-infrared SED of warm to hot dust emission offers promise for reducing the ambiguity of previous compact group studies. \citet[][hereafter J07]{kelsey07} have presented the first results from a Cycle~1 \spitzer\ IRAC (3.6--8.0\micron) and MIPS (24\micron) imaging survey of \total\ galaxies in 12 nearby HCGs. In brief, this work revealed trends between the evolutionary states of compact groups (determined from their dynamical and \HI\ masses), and their mid-infrared colors and luminosities. Galaxies in relatively gas-rich groups tend to have colors most indicative of star formation and AGN activity, and galaxies in gas-poor groups predominantly exhibit a narrow range of mid-infrared colors that are consistent with the light from quiescent stellar populations. The galaxies in this sample of 12 compact groups also occupy infrared color space in a distinctly different way than the population of galaxies in the {\it Spitzer} First Look Survey (FLS; e.g., \altcite{lacy04}), notably exhibiting a ``gap'' in their color distribution not found in the FLS sample. All of these results suggest that the environment of a compact group is intimately connected to the mid-infrared activity of the member galaxies. We present additional analysis of the nuclei of these galaxies, focusing in particular on developing diagnostics for identifying quantitative measures of mid-infrared activity with the ultimate goal of exploring the nature of mid-infrared emission, i.e., quiescent, star-forming, and/or accretion-powered. We also explore the connections between mid-infrared nuclear properties and other features of individual galaxies including morphology, optical spectral type, and the host group's evolutionary stage. Because of their proximity ($v<4500$\kms), we can spatially extract nuclear photometry for the HCG galaxies to reduce contamination from the extended galaxy that is inevitable in surveys of more distant galaxies. Throughout we assume a $\Lambda$-CDM cosmology with $\Omega_{\rm M}=0.3$, $\Omega_{\Lambda}=0.7$, and $H_0=70$\kms\,Mpc$^{-1}$ \citep[e.g.,][]{sperg07}. | Of the complete set of \total\ HCG galaxies presented in the {\em Spitzer} imaging survey of J07, we have identified 25 with red (\alphairac$<0.0$) mid-infrared continua. All eight known, spectroscopically identified star-forming galaxies (7a, 16c, 19c, 22c, 31ace, and 31b) are within this group. An additional three galaxies (16b, 61a, and 62b --- all optically identified AGNs) are apparently quiescent in the IRAC bands but show evidence for a 24\micron-excess above a stellar continuum, bringing the total number of 24\micron-active HCG galaxies to 61\% (as all of the mid-infrared active galaxies are also 24\micron-active). This is higher than the fraction of emission-line galaxies (including AGNs and star-forming galaxies) identified in optical spectroscopic surveys of HCGs \citep{coziol98a,coziol98b,shimada00}, and highlights the utility of multiwavelength studies to capture all activity. Though the mid-infrared analysis presented here has the advantage of being unambiguous in terms of identifying mid-infrared activity -- the distribution of HCG nuclei is clearly bimodal in terms of both \alphairac\ and \fracd\ -- distinguishing between accretion and star formation as the dominant source of mid-infrared power remains challenging. The evident trend in Figure~\ref{fig:alphalum24} between \alphairac\ and \ltwofour\ seen in mid-infrared active galaxy nuclei could be interpreted as an increasing contribution from star formation at higher \ltwofour\ and redder (more negative) values of \alphairac. Establishing this would be extremely powerful for interpreting photometric data of fainter objects in wide field surveys. However, the significant correlation between \frachd\ and \leight, which might similarly be read to imply a stronger AGN component (because of the greater hot dust contribution at 4.5\micron) at larger values of \leight, is not generally consistent with the known nuclear identifications from optical spectroscopy. This implies that simple measures of the steepness of the mid-infrared continuum are not diagnostic -- underscoring the difficulty of using mid-infrared photometry alone to categorize objects. Even with high signal-to-noise mid-infrared spectroscopy, the underlying source of mid-infrared power can remain ambiguous in some objects \citep[e.g.,][]{weedman05}. Furthermore, the strong stellar contamination at 3.6\micron\ in all of these galaxies means that color-color diagnostics that use the ratio of 3.6 to 4.5\micron\ fluxes to isolate AGNs \citep[e.g.,][]{stern05} likely miss a large fraction of them, particularly at lower infrared luminosities. Selection criteria with a wider dynamic range in color will be less sensitive to this effect \citep[e.g.,][]{lacy04}. However, as noted by J07, none of our sample would fulfill the AGN mid-infrared color-color selection criteria of \citet{stern05} or \citet{lacy04}, which were designed to target objects where the AGN dominates the mid-infrared SED. Our entire sample would also fail the power-law AGN selection criterion of \citet{donley07} that required a monotonic flux increase from 3.6 to 8.0\micron. A few optically identified LLAGNs -- 22a, 42a, and 62a -- show no evidence for any excess mid-infrared (including 24\micron) emission, indicating that warm to hot dust emission does not always accompany accretion onto a black hole. However, without a reservoir of cold material, an optically detectable AGN may be short-lived. In any case, such a system is unlikely to contribute significantly to the luminous energy budget of its host galaxy. Given their weak emission lines (which would be strongly diluted by host galaxy light at larger distances; \altcite{moran02}), these three galaxies may be identified with the optically dull but X-ray detected AGNs found in clusters \citep{martini06}. The evolved nature of their host groups and the X-ray detections of 42a and 62a in particular are consistent with identifying these systems as mini-clusters. In contrast, the less evolved compact groups may be less likely to host X-ray bright, optically quiet galaxies, as found in the loose group survey of \citet{shen+07}. At present, the lack of a complete suite of optical spectroscopy and high quality X-ray observations of our galaxy sample complicates definitively fitting the compact groups in context with the AGN surveys in the loose group and cluster environments done to date. The striking difference between the distribution of \alphairac\ values for the HCG and SINGS galaxies supports the role of environment in affecting HCG galaxies --- a conclusion consistent with the marked difference in the mid-infrared color-color distribution of HCG and FLS galaxies noted by J07. In particular, the strongly bimodal nature of HCG nuclear (and galaxy) \alphairac\ values as well as the correlation between group \HI\ content and 24\micron\ activity suggests that star formation in a compact group galaxy, likely induced by interactions, occurs in a burst that then exhausts its reservoir of cold gas, as suggested previously from far-infrared \citep{verdes98} and optical spectroscopic \citep{coziol98b} studies. The end result is a group dominated by mid-infrared quiescent galaxies. If local compact galaxy groups are analogous to the building blocks of clusters in the early Universe, this implies that the exhaustion of cold gas through enhanced star-formation occurs as a result of interactions (though not necessarily mergers) prior to cluster infall. Any AGN activity in those galaxies at that point would proceed in a low-luminosity, X-ray bright mode \citep{shen07}. However, Tully-Fisher \citep{Mendes2003} and fundamental plane studies \citep{delaRosa01} of HCG galaxies reveal that they are consistent with control galaxy populations. Therefore, there is no strong evidence for a significant effect of the compact group environment on the current states of galaxy morphology in such a significant way that would alter the end states of galaxy evolution. From a stellar population study of HCG ellipticals, \citet{delaRosa07} concluded that these galaxies showed evidence for truncated star formation in comparison to a (very small) sample of field ellipticals. While this may be the case, the mechanism for star formation truncation is unlikely to be extensive feedback from AGN activity. None of the optically known AGNs in any HCG has sufficient luminosity to clear a galaxy of its interstellar medium. The sole possible exception to this observation is 62a, where X-ray cavities in the intragroup medium suggest that a relativistic jet has ejected a sigificant amount of kinetic energy from a radiatively quiet AGN \citep[e.g.,][]{morita+06,gu+07}. However, this ongoing process likely postdates the epoch of star formation in this evolved group. | 7 | 10 | 0710.0372 |
0710 | 0710.0002_arXiv.txt | The spectral shape of the unresolved emission from different classes of gamma-ray emitters can be used to disentangle the contributions from these populations to the extragalactic gamma-ray background (EGRB). We present a calculation of the unabsorbed {\em spectral shape} of the unresolved blazar contribution to the EGRB starting from the spectral index distribution (SID) of resolved EGRET blazars derived through a maximum-likelihood analysis accounting for measurement errors. In addition, we explicitly calculate the {\em uncertainty} in this theoretically predicted spectral shape, which enters through the spectral index distribution parameters. We find that: (a) the unresolved blazar emission spectrum is only mildly convex, and thus, even if blazars are shown by GLAST to be a dominant contribution to the ERGB at lower energies, they may be insufficient to explain the EGRB at higher energies; (b) the theoretically predicted unresolved spectral shape involves significant uncertainties due to the limited constraints provided by EGRET data on the SID parameters, which are comparable to the statistical uncertainties of the observed EGRET EGRB at high energies; (c) the increased number statistics which will be provided by GLAST will be sufficient to reduce this uncertainty by at least a factor of three. | The isotropic, and presumably, extragalactic gamma-ray background emission (EGRB) detected by the Energetic Gamma-ray Experiment Telescope (EGRET) aboard the {\it Compton Gamma-ray Observatory} (Sreekumar et al.\ 1998) is one of the most important observational constraints on known or theorized populations of faint, unresolved gamma-ray emitters. With the imminent launch of the {\it Gamma-ray Large Area Space Telescope} (GLAST), which is expected to represent an unprecedented leap in observational capabilities in GeV energies, the timing is especially opportune to consider the information content of the diffuse background and methods for maximizing the scientific return from its study. One of the primary challenges in using the EGRB to constrain properties of extragalactic gamma-ray emitters and exotic physics is disentangling the convolved contributions of guaranteed participating populations. Estimates of the levels of the collective unresolved emission even from established classes of extragalactic sources (such as blazars and normal galaxies) involve significant uncertainties and are at the order-of-magnitude level at best (e.g., Padovani et al.\ 1993; Stecker \& Salamon 1996; Kazanas \& Perlman 1997; Mukherjee \& Chiang 1999; M\"ucke \& Pohl 2000; Dermer 2006; Lichti et al.\ 1978; Pavlidou \& Fields 2002). A very promising approach for the study of the EGRB and its components is through the use of spectral shape information. Let us consider the optimal case where the expected spectral shapes of the unresolved emission from known classes of gamma-ray sources can be confidently predicted. In this case, a series of conclusions can be drawn regarding the potential contributions of these classes to the EGRB even without detailed calculations of the {\em magnitudes} of their collective emission. For example, in comparing the spectral shape of the spectrum due to a particular class with that of the EGRB, one can identify whether this class could, in principle, comprise most of the EGRB or require the existence of contributions from other classes (Stecker \& Salamon 1996a,b; Strong et al. 2004; Pavlidou \& Fields 2002)\footnote{Note that the method of spectral comparison can only be used to {\em reject} a population from being the sole source of the EGRB; spectral consistency does not constitute in itself proof of the importance of a class of objects as an EGRB contributor, since the overall normalization of the emission may, in fact, be low depending on the gamma-ray luminosity function of the population.}. Potentially identifiable spectral features could be predicted and searched for in GLAST data (e.g., Pavlidou \& Fields 2002). Finally, spectral information could be used to ultimately disentangle different components and contributions (as in e.g. de Boer et al.\ 2004 for the case of the diffuse emission from the Milky Way). An additional attractive feature of such calculations is that the associated uncertainties are largely independent of those entering the calculations of the overall unresolved emission flux. As blazars are the most populated class of gamma-ray emitters, unresolved blazars are guaranteed to contribute significantly, if not dominantly, to the EGRB. Thus, it is especially important to understand the expected spectral shape of their collective unresolved emission and the uncertainties involved in its calculation. Individual blazars have been measured to have power-law spectra in the EGRET range, $F_E \propto E^{-\alpha}$, with spectral index $\alpha$ ranging approximately from 1.5 to 3. The unresolved emission from a collection of power-law emitters with variable spectral indices\footnote{Statistically significant spread in the observed and intrinsic spectral index of blazars has also been confirmed in other energy bands (see e.g.\ Shen et al.\ 2006 for the case of X-ray emission).} has, invariably, a convex spectral shape (Brecher \& Burbidge 1972; Stecker \& Salamon 1996; Pohl et al.\ 1997; Pavlidou et al.\ 2007). The exact shape of the unresolved spectrum depends on the spectral index distribution (SID) of gamma-ray loud blazars. Recognizing that this is the case, Stecker \& Salamon (1996a) explicitly reconstructed a spectrum from the observed SID of blazars in the 2nd EGRET catalog (Thompson et al.\ 1995), deriving a spectrum that was indeed significantly convex. However, measurement errors in individual spectral indices smear the SID and exaggerate the curvature of the spectrum (Pohl et al.\ 1997). Recently, in Venters \& Pavlidou 2007 (hereafter VP07), we have applied a maximum-likelihood analysis to recover the intrinsic spectral index distribution (ISID) of gamma-ray loud blazars from EGRET observations. We found that (1) the maximum-likelihood ISID is appreciably narrower than the observed SID, so the best-guess spectrum is likely to have only a mild curvature; (2) BL Lacs and flat spectrum radio quasars (FSRQs) are likely to be spectrally distinct populations with spectrally distinct contributions to the EGRB; (3) there is no evidence for a systematic shift of spectral variability with flaring implying that although variability may be important in the level of the contribution from blazars, ignoring variability effects in spectral shape studies is likely to be a good approximation. Here, we use the ISIDs derived in VP07 to calculate the spectral shape of the blazar contribution to the EGRB. We examine the sensitivity of the shape to the exact values of the ISID parameters and report on the range of possible shapes given our uncertainties in the determination of these parameters. We also investigate how the spectral shapes of the BL Lac and FSRQ contributions may differ. Finally, we predict how our understanding of the spectral shape of the unresolved blazar emission will improve after GLAST observations become available. | \label{Disc} We have calculated the expected spectral shape of the unresolved gamma-ray emission from blazars under the assumptions that the ISIDs of blazars do not evolve with redshift and are independent of blazar luminosity and flaring state. We have also explicitly calculated the $1\sigma$ uncertainty in the spectral shape entering through the limited constraints on the blazar ISIDs derived from EGRET data. Finally, we have predicted by how much these uncertainties will be reduced if GLAST observations are used to determine the blazar ISIDs. The unresolved emission spectral shape can be used as an indicator of the potential importance of a given population's contribution to the EGRB, and it constrains the maximal contribution at high energies relative to that at low energies. If the curvature tentatively seen in the observed EGRB is real, then a population with little such curvature in its unresolved spectrum (such as the FSRQs) will not be the dominant contributor to the EGRB at high energies even if it is dominant at low energies. The unresolved BL Lac spectrum does seem to be more convex than that of the FSRQs, but the level of uncertainty in the spectral shape is high in this case because only a few BL Lacs were detected by EGRET. However, GLAST observations will dramatically improve our understanding of the blazar ISIDs and the associated unresolved spectral shapes allowing us to use spectral shape information to calculate the minimal additional contribution required from other classes of sources to explain the observed EGRB spectrum. It should be stressed that here we have only calculated {\em unabsorbed} spectra. At energies higher than 10 ${\rm \, GeV}$, gamma-ray absorption through interactions (pair production) with the extragalactic background light becomes important (e.g., Salamon \& Stecker 1998), and the spectral shape of any contribution to the EGRB will be accordingly changed. We will return to this effect in a future publication. | 7 | 10 | 0710.0002 |
0710 | 0710.2225_arXiv.txt | We have determined the mineralogical composition of dust in the Broad Absorption Line (BAL) quasar \pgbal\ using mid-infrared spectroscopy obtained with the {\em Spitzer Space Telescope}. From spectral fitting of the solid state features, we find evidence for Mg-rich amorphous silicates with olivine stoichiometry, as well as the first detection of corundum (Al$_2$O$_3$) and periclase (MgO) in quasars. This mixed composition provides the first direct evidence for a clumpy density structure of the grain forming region. The silicates in total encompass ($56.5 \pm 1.4$)\% of the identified dust mass, while corundum takes up ($38 \pm 3$) wt.\%. Depending on the choice of continuum, a range of mass fractions is observed for periclase ranging from ($2.7 \pm 1.7$)\% in the most conservative case to ($9 \pm 2$)\% in a less constrained continuum. In addition, we identify a feature at 11.2 $\mu$m as the crystalline silicate forsterite, with only a minor contribution from polycyclic aromatic hydrocarbons. The ($5 \pm 3$)\% crystalline silicate fraction requires high temperatures such as those found in the immediate quasar environment in order to counteract rapid destruction from cosmic rays. | \label{sec:intro} In the local Universe, Asymptotic Giant Branch (AGB) stars are believed to dominate dust production \citep[e.g.,][]{D_03_dust}. AGB stars are a late phase of stellar evolution for stars with initial masses of $M$=1--8\msun, which develop outflows where the low temperatures and high densities become optimal for dust condensation. However, a significant amount of dust originating from AGB stars does not build up within $\sim1$~Gyr of the birth of the first generation of low-mass stars \citep{ME_03_earlydust}. Nevertheless, large quantities of dust are clearly present at these early times: quasar host galaxies at $z\sim6$ show evidence for 10$^{8-9}$~\msun\ of dust heated by star formation in their host galaxies as seen in submm and far-infrared emission \citep{PriddeyEtal2003,BCB_06_dust}. Furthermore, extinction curves for a $z=6.2$ quasar and GRB~050904 ($z=6.29$) are not consistent with those locally \citep{maio04,stratta07}. Another mechanism for dust production is therefore required, and \citet{elvis-dust} explored the possibility that quasar winds might provide environments suitable for efficient dust formation. They determined that the temperatures and pressures in these regions could reach the values found in cool, dust-producing stars. The quasar winds are predicted to produce dust masses up to $10^7$ \msun\ \citep{elvis-dust}, implying that the quasar wind may in some cases only account for part of the dust production, and supernovae have been suggested as an alternative source of dust in high-z galaxies \citep{SEB_06_SNdust}. Quasar outflows are most obviously manifested in Broad Absorption Line (BAL) quasars. This population, approximately 20\% of optically selected type~1 quasar samples \citep[e.g.,][]{HewFol2003,reich+03a}, is notable for broad, blueshifted absorption evident in common ultraviolet resonance transitions such as \CIV, Ly$\alpha$, and \OVI. These P~Cygni-type features arise because the observer is looking through an outflowing wind. \balqs\ are thus a natural population to consider when investigating the grain properties of dust in the quasar environment. In this letter, we present a detailed analysis of the mid-infrared spectrum of the luminous \balq, \pgbal, in order to determine its dust composition. \pgbal, an {\em IRAS} source, is mid-infrared bright, and one of the most luminous low-redshift ($z=0.466$) Bright Quasar Survey objects with $M_{\rm V}=-26.9$ (\hnot$=70$\hunits, $\Omega_{\rm M} = 0.3$, and $\Omega_{\Lambda} = 0.7$ are assumed throughout). \hst\ spectra revealed broad, shallow \CIV\ absorption, and it has been well-studied in the UV and X-ray \citep[e.g.,][]{gall+04}. | For the first time, the composition of the dust in a quasar has been determined, albeit limited to the components with resonances in the infrared. While the origin of the dust could lie in stellar ejecta, the dust properties are consistent with their formation in the quasar wind itself. The dust in \pgbal\ clearly bears similarity to dust in other astrophysical environments, i.e.~in the properties of the amorphous silicates, but the presence of large amounts of MgO and Al$_2$O$_3$ sets it apart from the composition of interstellar dust in the local universe. Assuming the dust is formed in the quasar wind, the co-existence of the highly refractory corundum, the crystalline silicates and and the less refractory MgO and amorphous silicates indicates that the wind of \pgbal\ has an inhomogeneous temperature and density structure. | 7 | 10 | 0710.2225 |
0710 | 0710.0827_arXiv.txt | {Multiple systems are the product of protostellar core fragmentation. Studying their statistical properties in young stellar populations therefore probes the physical processes at play during star formation.} {Our project endeavors to obtain a robust view of multiplicity among embedded Class\,I and Flat Spectrum protostars in a wide array of nearby molecular clouds to disentangle ``universal'' from cloud-dependent processes.} {We have used near-infrared adaptive optics observations at the VLT through the $H$, $K_s$ and $L'$ filters to search for tight companions to 45 Class\,I and Flat Spectrum protostars located in 4 different molecular clouds (Taurus-Auriga, Ophiuchus, Serpens and L1641 in Orion). We complemented these observations with published high-resolution surveys of 13 additional objects in Taurus and Ophiuchus.} {We found multiplicity rates of 32$\pm$6\% and 47$\pm$8\% over the 45--1400\,AU and 14--1400\,AU separation ranges, respectively. These rates are in excellent agreement with those previously found among T\,Tauri stars in Taurus and Ophiuchus, and represent an excess of a factor $\sim$1.7 over the multiplicity rate of solar-type field stars. We found no non-hierarchical triple systems, nor any quadruple or higher-order systems. No significant cloud-to-cloud difference has been found, except for the fact that all companions to low-mass Orion protostars are found within 100\,AU of their primaries whereas companions found in other clouds span the whole range probed here.} {Based on this survey, we conclude that core fragmentation always yields a high initial multiplicity rate, even in giant molecular clouds such as the Orion cloud or in clustered stellar populations as in Serpens, in contrast with predictions of numerical simulations. The lower multiplicity rate observed in clustered Class\,II and Class\,III populations can be accounted for by a universal set of properties for young systems and subsequent ejections through close encounters with unrelated cluster members.} | The prevalence of binary and multiple systems among stellar populations in our Galaxy is generally understood as a consequence of the natural tendency of prestellar cores for fragmentation during or immediately after their free-fall collapse (see Tohline 2002 for a detailed review). Numerical simulations have long predicted that this fragmentation, hence the frequency and properties of multiple systems, is strongly dependent on the initial conditions reigning in the core (e.g., Bonnell et al. 1992; Durisen \& Sterzik 1994; Boss 2002; Goodwin et al. 2004b). Millimeter observations over the last two decades have provided a detailed view of the initial conditions of star formation. For instance, the basic properties of individual prestellar cores, outer radius and central density, differ significantly from one molecular cloud to another (Motte \& Andr\'e 2001). It is also likely that the temperature in prestellar cores vary from cloud to cloud depending on the strength of the interstellar radiation field (e.g., Stamatellos et al. 2007). Finally, there is now good evidence that protostellar collapse is generally more violent in a cluster-forming environment than in isolated dense cores, e.g., induced by strong external disturbances as opposed to spontaneous or self-initiated (Andr\'e et al. 2003; Belloche et al. 2002, 2006). Considering these differences, one may therefore expect to observe substantial differences in the properties of multiple systems in independent stellar populations. Early high-angular resolution surveys of Myr-old T\,Tauri stars in nearby T associations, such as Taurus and Ophiuchus (Ghez et al. 1993; Leinert et al. 1993; Reipurth \& Zinnecker 1993) revealed a significant excess over the multiplicity rate among field stars (Duquennoy \& Mayor 1991, hereafter DM91; Fischer \& Marcy 1992). It was then rapidly discovered that clustered populations of equally young stars, such as the Orion Nebula Cluster (ONC), do not possess such a high multiplicity rate and rather resemble field stars from this point of view (Padgett et al. 1997; Petr et al. 1998; Duch\^ene et al. 1999). While this could be evidence that supports the environment-dependent fragmentation of prestellar cores, alternative explanations cannot be excluded. In particular, dynamical interactions with unrelated cloud members can disrupt most wide companions in less than 1\,Myr in the densest clusters (Kroupa 1995). In other words, T\,Tauri multiple systems have had sufficient time to evolve since their formation, so that their observed properties may not be considered as directly representative of core fragmentation. To circumvent this problem, observations of less evolved, embedded, young stellar objects (YSOs) are required in order to determine the pristine properties of multiple systems. Radio interferometric studies of embedded Class\,0 and Class\,I sources have hinted at a high overall multiplicity rate, comparable to that observed among T\,Tauri stars in non-clustered populations (Looney et al. 2000; Reipurth et al. 2002; 2004). Non uniform and limited samples, as well as the fact that not all YSOs emits strongly at centimeter wavelengths, prevent conclusive comparisons at this point, however. In parallel to this effort, Haisch et al. (2002, 2004) and Duch\^ene et al. (2004, hereafter D04) conducted direct near-infrared imaging surveys of Class\,I protostars. These surveys also found a high multiplicity rate, consistent with that of somewhat more evolved T\,Tauri stars in the same clouds. D04 found marginal evidence that the multiplicity rate of these sources decreases on a timescale of $\sim10^5$\,yr, possibly a result of the internal or external dynamical ejection of wide companions. No strong evidence was found for a variation of the multiplicity rate of Class\,I sources from one cloud to another, in part because of small sample sizes and of the limited range of projected separations probed by these surveys. Whether fragmentation does indeed depend on environmental conditions remains to date a theoretical/numerical predictions that has not yet been confirmed observationally. In particular, the outcome of core fragmentation in a very rich molecular cloud such as Orion remains unknown. To extend the analysis of D04, we have undertaken a high-angular resolution survey of embedded Class\,I protostars sampling the Serpens and Orion molecular clouds in addition to Taurus and Ophiuchus. This paper is organized as follows: we present our sample and observations in Section\,\ref{sect:obs} and the results in Section\,\ref{sect:results}. In Section\,\ref{sect:discus}, we analyze these results in view of other multiplicity surveys and of predictions of star formation theories. We summarize our findings in Section\,\ref{sect:concl}. | \label{sect:concl} We have used diffraction-limited imaging from 1.6 to 3.7\,$\mu$m with the 8m-VLT adaptive optics system to search for tight companions around 45 embedded Class\,I and FS protostars in the Taurus, Ophiuchus, Serpens and L1641 (Orion\,A) molecular clouds. We complement our analysis with published high-resolution surveys of similar objects in Taurus and Ophiuchus to build a sample of 58 Class\,I and FS targets. We derive an average multiplicity rate of 32$\pm$6\% over the separation 45--1400\,AU. In the Taurus and Ophiuchus clouds, the closest clouds in our sample, we further derive a multiplicity rate of 47$\pm$8\% over the separation range 14--1400\,AU. These rates are a factor of $\sim1.7$ higher than those derived for nearby solar-type field stars, extending the multiplicity excess found among several populations of T\,Tauri star to even younger ages. Most importantly, we find that the embedded protostars in L1641 show a multiplicity rate similar to that in other clouds, indicating for the first time that a high multiplicity rate is achieved after core fragmentation in all types of nearby molecular clouds, including giant molecular clouds such as Orion\,A, which also hosts the dense ONC cluster. We also find a high multiplicity rate in Serpens, the densest cluster in our survey. Our results support the view that core fragmentation results in a multiplicity rate for wide companions that does not depend on the initial conditions reigning in the cores, as opposed to predictions of most numerical simulations. Rather, our findings support a scenario in which all YSO populations start with a similar set of multiplicity properties and only evolve as a consequence of disruptive system-system interactions prior to dilution of the clusters in the field. Follow-up multiplicity surveys of the Class\,II and Class\,III populations of the Serpens and L1641 clouds would provide key empirical verifications of this scenario. We found 6 triple systems, all of them hierarchical, and higher-order systems are very rare among Class\,I/FS sources within the separation ranges studied here. It is unlikely that many companions will be ejected as a result of internal decay of unstable systems past the Class\,I phase. Finally, we also find a possible trend for the embedded Orion multiple systems to have different orbital properties than those in other clouds, namely systematically tighter projected separations, a hint that environmental conditions may impact on the properties of protobinaries. | 7 | 10 | 0710.0827 |
0710 | 0710.5891_arXiv.txt | { Spectral differential imaging is an increasingly used technique for ground-based direct imaging searches for brown dwarf and planetary mass companions to stars. The technique takes advantage of absorption features that exist in these cool objects, but not in stars, and is normally implemented through simultaneous narrow-band imagers in 2 to 4 adjacent channels. However, by instead using an integral field unit, different spectral features could be used depending on the actual spectrum, potentially leading to greater flexibility and stronger detection limits. In this paper, we present the results of a test of spectral differential imaging using the SINFONI integral field unit at the VLT to study the nearby active star L449-1. No convincing companion candidates are found. We find that the method provides a $3 \sigma$ contrast limit of 7.5 mag at 0.35", which is about 1.5 mag lower than for NACO-SDI at the same telescope, using the same integration time. We discuss the reasons for this, and the implications. In addition, we use the SINFONI data to constrain the spectral type in the NIR for L 449-1, and find a result between M3.0 and M4.0, in close agreement with a previous classification in the visual range. | In recent years, developments in techniques and instrumentation have led to a strongly increased capacity for high-contrast imaging of substellar companions at small angular separations to stars, from the ground. In particular, the use of high-order adaptive optics (AO) in combination with various differential imaging techniques provides a sensitivity for companions that are $10^4$-$10^5$ times fainter than the primary at a separation of 0.5"-1" at near-infrared wavelengths (see Janson et al. 2007 and Biller et al. 2007). According to theoretical models (see e.g. Baraffe et al. 2003), this corresponds to objects of a few times the mass of Jupiter around young stars. A particularly efficient technique for such observations is simultaneous spectral differential imaging (SDI, see e.g. Rosenthal et al. 1996 and Racine et al. 1999), in which the flux from a star is observed in two narrow, adjacent wavelength bands simultaneously. The bands are chosen such that one is inside and one outside of an absorption feature which arises uniquely in cool atmospheres. One example of such a feature is the methane band at 1.6 $\mu$m, which only occurs in dwarfs of spectral type T or later, whereas a stellar spectrum is flat in this range. By subtracting one image from the other, the stellar point spread function (PSF) can be largely subtracted out, since it is approximately flat around 1.6 $\mu$m. Along with the main PSF, most of the random PSF substructure (speckle noise) is subtracted out as well, since the images are simultaneous and hence the atmospheric speckle pattern is the same in both images. NACO-SDI at the VLT is an excellent instrument for the purpose of SDI observation (see Lenzen et al. 2004). NACO-SDI images the same field in three different narrow bands around the 1.6 $\mu$m methane feature simultaneously. It is designed to minimize non-common path aberrations, and has been shown to be capable of achieving a contrast of more than 13 mag between star and companion at 1" angular separation at the 3$\sigma$ level (see Janson et al. 2007), without the use of a coronagraph. One limitation of the NACO-SDI instrument is that it can only attain information about the methane feature in question. It has been suggested (Berton et al. 2006) that if $N > 1$ different absorption features are used at once for SDI purposes, the achievable contrast would increase by a factor $N^{1/2}$ for the same integration time, simply due to the more efficient use of photons. Hence, in principle, an instrument with the capacity to do multi-feature SDI (MSDI) would be preferable to NACO-SDI, if the instruments perform equally well in all other respects (in terms of differential aberrations, etc.). SINFONI is an integral field spectroscopy instrument with AO capability at the VLT (see Eisenhauer et al. 2003 and Bonnet et al. 2004). It uses an image slicer to divide the FOV into pseudo-slits that are individually dispersed on a grating, hence retaining both the spatial and spectral information of the incoming light. The output data can be used to construct a data cube that stores spatial information along two axes, and spectral information along the third. In this way, SINFONI can be used as an SDI or MSDI instrument, since the cube contains simultaneous narrow-band images over a large range of wavelengths. We have acquired SINFONI data to test its capacity for SDI. In the following sections, we describe the observations and data reduction, and compare the results with NACO-SDI. We also compare the results to a spectral deconvolution (SD) scheme applied to another set of SINFONI data by Thatte et al. (2007). We dicuss the advantages and disadvantages of using SINFONI for high-contrast imaging and characterization of exoplanet and brown dwarf companions to stars. Finally, we analyze the non-differential collapsed images, and constrain the spectral type of L 449-1 from H- and K-band spectroscopy. | We have investigated the potential of the SDI technique with SINFONI for detecting cool substellar companions to stars. It was discovered that the performance is considerably worse than for NACO-SDI under very similar circumstances. In addition, the error is partly static, leading to a poor increase of performance with increasing integration time. The quasi-static nature of the noise, along with the fact that it is constant over most of the image space, implies that non-common path aberrations and errors from the slicing and reconstruction of the image dominate the residual noise. These results, along with practical considerations, clearly lead to the suggestion that ``blind'' differential imaging searches for cool companions (i.e., when no a priori information is available) are presently best performed with specialized instruments such as NACO-SDI, whereas follow-up observations, where more qualitative information is desired, are better suited for integral field units such as SINFONI. The results are highly relevant for the design of the next-generation instruments for planet detection. While the collapsed broad-band images of L 449-1 imply the existence of a low-contrast candidate companion, the absence of a clear spectral signature associated with this feature suggests that it is a PSF artefact rather than a physical object. The spectral type of L449-1, as determined from the H- and K-band spectra, was found to be in the range of M3.0-M4.0, in close agreement with previous results (M4, Scholz et al. 2005) in the visual range. | 7 | 10 | 0710.5891 |
0710 | 0710.0961_arXiv.txt | The transient X-ray binary pulsar A0535+262 was observed with $Suzaku$ on 2005 September 14 when the source was in the declining phase of the August-September minor outburst. The $\sim$103 s X-ray pulse profile was strongly energy dependent, a double peaked profile at soft X-ray energy band ($<$ 3 keV) and a single peaked smooth profile at hard X-rays. The width of the primary dip is found to be increasing with energy. The broad-band energy spectrum of the pulsar is well described with a Negative and Positive power-law with EXponential (NPEX) continuum model along with a blackbody component for soft excess. A weak iron K$_\alpha$ emission line with an equivalent width $\sim$ 25 eV was detected in the source spectrum. The blackbody component is found to be pulsating over the pulse phase implying the accretion column and/or the inner edge of the accretion disk may be the possible emission site of the soft excess in A0535+262. The higher value of the column density is believed to be the cause of the secondary dip at the soft X-ray energy band. The iron line equivalent width is found to be constant (within errors) over the pulse phase. However, a sinusoidal type of flux variation of iron emission line, in phase with the hard X-ray flux suggests that the inner accretion disk is the possible emission region of the iron fluorescence line. | Be X-ray binaries consist of a neutron star in an eccentric orbit around a Be star companion. The orbit of the Be X-ray binaries is generally wide and eccentric with orbital periods in the range of 16 days to 400 days (Coe 2000). Mass transfer from the Be companion to the neutron star takes place through the circumstellar disk. When the neutron star passes through the disk or during the periastron passage, it shows strong outbursts with an increase in X-ray luminosity by a factor $\geq$ 100 (Negueruela 1998). A0535+262 is a 103 s Be/X-ray binary pulsar discovered by $Ariel~V$ during a large outburst in 1975 (Coe et al. 1975). The binary companion HDE~245770 is an O9.7-B0 IIIe star in a relatively wide eccentric orbit ($e$ = 0.47) with orbital period of $\sim$ 111 days and at a distance of $\sim$ 2 kpc (Finger et al. 1996; Steele et al. 1998). The pulsar shows regular outbursts with the orbital periodicity. Occasional giant X-ray outbursts are also observed when the object becomes even brighter than the Crab. The pulsar shows three typical intensity states, such as quiescence with flux level of below 10 mCrab, normal outbursts with flux level in the range 10 mCrab to 1 Crab, and giant outbursts during which the object becomes the brightest X-ray source in the sky with the flux level of several Crab (Kendziorra et al. 1994). 103 s pulsations were detected during past X-ray observations of A0535+262 in quiescence, outburst, and giant outbursts. The pulse profile was single peaked in quiescence (in 1-10 keV range; Mukherjee \& Paul 2005, 3-20 keV; Negueruela et al. 2000), double peaked with a clearly asymmetric ''main'' and a more symmetric ''secondary'' pulse during the X-ray outbursts (Mihara 1995; Kretschmar et al. 1996, Maisack et al. 1997). The X-ray spectrum of the pulsar has been studied at soft and hard X-rays with various instruments at different luminosity levels. The spectrum of the object, during outbursts, shows cyclotron resonant scattering features at higher energies than that of the other pulsars. Two harmonic features at around 50 keV and 100 keV were detected in its 1989 outburst with the HEXE/TTM instrument on Mir/Kvant (Kendziorra et al. 1994). The CGRO/OSSE observations of 1994 outburst of the pulsar showed a significant absorption feature at 110 keV (Grove et al. 1995). These detections did not resolve whether the magnetic field of the pulsar is $\sim$ 5 $\times$ 10$^{12}$ G (when the fundamental occurs at 55 keV) or $\sim$10$^{13}$ G (for the 110 keV fundamental). The most recent major X-ray outburst was detected in 2005 May/June with the BAT instrument on Swift when the 15-195 keV count rate was greater than 3 times that of the Crab Nebula (Tueller et al. 2005). Following the detection of the recent outburst, the pulsar was observed by INTEGRAL, RXTE and the recently launched $Suzaku$. A cyclotron resonance feature at $\sim$45 keV was detected in the Hard X-ray Detector (HXD) spectrum of $Suzaku$ with estimated magnetic field of the pulsar as $\sim$ 4 $\times$ 10$^{12}$ Gauss (Terada et al. 2006). The detection of the absorption feature at $\sim$45 keV and its first harmonic at $\sim$100 keV is reported from the INTEGRAL and RXTE observations of the pulsar during the 2005 August/September outburst (Caballero et al. 2007). Using the same Suzaku observation used for the analysis of the cyclotron resonance feature (Terada et al. 2006), we study the broad-band spectral properties of the pulsar in the present paper. | Apart from two RXTE observations and three BeppoSAX observations during the quiescence, all other observations reported before 2005 August were made during the giant outbursts when the source flux was about equal to or more than that of the Crab nebula. The RXTE and INTEGRAL observations of the pulsar, during the 2005 August-September minor outburst, confirmed the earlier discovery of a cyclotron feature at $\sim$45 keV along with the detection of its first harmonic at $\sim$100 keV (Caballero et al. 2007). Suzaku observation of A0535+262 during the declining phase of the 2005 August-September normal outburst provided new data on spectral properties of the pulsar up to $\sim$100 keV and how they might evolve throughout the outburst. \subsection{Pulse Profile} During the present observation of the minor outburst, the pulse profile of A0535+262 is different from that in the quiescence or during the earlier reported outbursts. The pulse profile in the hard energy band is known to show energy and luminosity dependence in outbursts, which is described in Bildsten et al. (1997) based on the BATSE observations. The profile tends to be double peaked when the luminosity is high, eg.\ $>\!10^{37}$ erg s$^{-1}$ (20--100 keV), whereas it becomes a single peak at $4\times10^{36}$ erg s$^{-1}$. The current $Suzaku$ observation had even lower luminosity ($2\times10^{35}$ erg s$^{-1}$ in 20--100 keV band), and showed almost sinusoidal profile above 25~keV\@. The $Suzaku$ hard band observation found a single peaked profile, as was seen by BATSE. Hard X-ray pulse profiles of A0535+262, obtained from the RXTE and INTEGRAL observations of the same 2005 outburst with a peak luminosity of $\sim$10$^{37}$ erg s$^{-1}$, appear to be double-peaked (Caballero et al. 2007). The two peaks with a shallow valley constitute a trapezoid-like profile above 35 keV separated by a wide trough. Though the observations were carried out during the peak and the decay of the outburst, the added pulse profile may be dominated by the profile during the peak of the outburst. Therefore, the profile presented by Caballero et al. (2007) probably reflect that at the higher luminosity observations (at the peak of the outburst). The double-peaked hard X-ray profiles at a source luminosity of $\sim$10$^{37}$ erg s$^{-1}$ agrees with the luminosity dependence of pulse profiles observed with BATSE. The profiles obtained from Suzaku observation, however, appear to be different from that of the RXTE and INTEGRAL observations. This is interpreted due to the luminosity dependent effect as the Suzaku observation was at a luminosity level of two orders of magnitude less compared to the peak luminosity. On the other hand, the Suzaku pulse profile in the soft energy band ($<\!10$ keV) is found to be similar to that obtained from the $Ginga$ observation of the 1989 outburst (Mihara 1995). The dip-like structure, prominent only at soft X-rays is also recognized in the $Ginga$ profile. Because the $Ginga$ observation was made during the outburst when the source was $\sim$80 times brighter than the present observation, the profile may be interpreted that A0535+262 in outbursts generally contains a dip-like structure at soft X-rays as seen in 2005 $Suzaku$ observation. This type of structure is also seen in the RXTE observation of the pulsar during the same minor outburst in 2005 August/September (Caballero et al. 2007). This dip-like structure is absent in the pulse profiles when the pulsar was in quiescence (Mukherjee \& Paul 2005). As other observations of A0535+262 in giant outbursts focused on the hard X-ray energy ranges, the dip-like structure could have been missed in the pulse profiles. \subsection{Phase-Averaged Spectroscopy} The X-ray spectrum of A0535+262 has been described by different continuum models at different energy ranges. Hard X-ray observations during outbursts, meant for the study of cyclotron absorption features and hence the pulsar magnetic field, were fitted by different continuum models such as optically thin thermal bremsstrahlung, power-law, power-law with an exponential cut-off, Wien's law or different combinations of above components (dal Fiume et al.\ 1988; Kendziorra et al.\ 1994; Grove et al.\ 1995), whereas the 2--37 keV spectrum, obtained from the $Ginga$ observation of the 1989 outburst, was well fitted with NPEX continuum model (Mihara 1995). Current analysis of the $Suzaku$ data showed that the NPEX continuum model fits well to the source spectrum. Selection of an appropriate continuum model may be crucial to investigate the broad features in the spectrum, such as the soft excess represented by a blackbody component. The spectral fitting to the 0.3--10~keV BeppoSAX spectra, during quiescence, yielded a blackbody component of temperature $\sim$1.3~keV along with a power-law component of index $\sim$1.8 (Mukherjee \& Paul 2005). On the other hand, the Suzaku data when fitted with NPEX continuum model yielded a much lower temperature of $\sim$0.16 keV, which is common for most of the X-ray pulsars. We suspect that the difference arises from the different selection of the continuum model. In fact, when the Suzaku wide band spectra were fitted by a power-law with an exponential cutoff, we did indeed obtain a blackbody temperature of 1.36 keV. This demonstrates the importance of the continuum selection to quantify the blackbody component. However, the selection criteria for the continuum choice have yet to be rigorously defined. Most of the transient Be/X-ray binary pulsars undergo periodic outbursts due to the enhanced mass accretion when the neutron star passes through the dense regions of the circumstellar disk. During the passage, there is every possibility of the increase in the absorption column density. In case of A0535+262, we found that the absorption column density was higher than the Galactic column density towards A0535+262 (5.9 $\times$ 10$^{21}$ atoms cm$^{-2}$) during the minor outburst in 2005 August-September. Note that the $Suzaku$ observation was made near the apastron at the orbital phase of 0.42--0.43. From BeppoSAX observations of the pulsar in quiescence (orbital phases 0.77, 0.05, 0.42), though, the estimated value of the absorption column density was found to be similar to that of the Galactic column density (Mukherjee \& Paul 2005). During the 1998 RXTE/PCA observations (in quiescence, orbital phases of 0.0 and 0.77), the value of the absorption column density was found to be a factor of about 10 higher than that of the Galactic column density (Negueruela et al. 2000). During the $Ginga$ observation of the pulsar in the 1989 outburst (orbital phase of 0.96), the estimated absorption column density was found to be 4.8$^{+1.7}_{-1.2}$ $\times$ 10$^{21}$ atoms cm$^{-2}$ (Mihara 1995), which is compatible with the Galactic value. These are only three instances where the absorption column density of the pulsar is reported so far. Although the available data are scarce, the absorption column density does not show clear correlation with the source intensity or the orbital phase. In this sense, it is note worthy that the observations of the large absorption column with RXTE/PCA were made when the circumstellar disk around the Be companion (Be disk) was absent (Negueruela et al. 2000). Haigh et al. (2004) suggest that the X-ray activity of the pulsar is correlated with the change of the truncation radius of the Be disk. The truncation of the disk by the tidal interaction with the neutron star, which defines the truncation radius, could explain the observed X-ray outbursts in Be/X-ray binary systems. When the truncation radius reduces, the material from the outer portions of the disk are expelled and may be found elsewhere in the binary system. This may trigger the giant X-ray outburst. When the Be disk disappears, as the case of RXTE/PCA observations, all the disk material is considered to be accreted and/or ejected throughout the binary system. The presence of the expelled material in the line of sight may have caused a moderately large absorption column. Haigh et al. (2004) argues that the reduction of the Be disk continues for several orbital periods before/after the giant outburst. Because the Suzaku observation was made about an orbital period after the giant outburst in May/June 2005, the reduced Be disk may have caused moderately large absorption column. Broad-band spectroscopy of A0535+262 shows the presence of a weak iron K$_\alpha$ emission line with $\sim$25~eV equivalent width. Other than the 1989 $Ginga$ observation and recent RXTE and INTEGRAL observations of the same minor outburst, none of the previous observations of the pulsar could detect the weak emission line at $\sim$6.4~keV\@. The emission line is interpreted as due to the fluorescent line from the cold ambient matter in the vicinity of the neutron star. If the cold matter is distributed spherically symmetric around the neutron star, the excess column density over the Galactic value, $\sim\!5\times10^{21}$ atoms cm$^{-2}$, would produce an iron emission line with an equivalent width of only a few eV (Makishima 1986). The larger equivalent width we detected may be interpreted as the cold matter having an asymmetric distribution, and more matter is distributed outside the line of sight. The fluorescence of the atmosphere of the binary companion could be a possible source of the iron emission line. However, we consider that it does not have a major contribution to the observed equivalent width, because the solid angle subtended by the companion, which is estimated as $\Omega/4\pi$ $\leq$ 10$^{-4}$ at apastron, is very small due to the long-orbital period of A0535+262. Instead, we conjecture that the accretion disk is the probable reprocessing site to produce the iron fluorescent line. Because the pulse phase with higher line flux covers more than a half of the pulse period, the reprocessing site should subtend a large solid angle against the neutron star. The accretion disk conforms to this condition. \subsection{Pulse Phase Resolved Spectroscopy} Pulse phase resolved spectroscopy of A0535+262 showed that the blackbody component, used to describe the soft excess, is pulsating as seen in some other accreting X-ray pulsars, such as LMC~X-4, SMC~X-1 etc. (Naik \& Paul 2004a, Naik \& Paul 2004b, Paul et al.\ 2002 and references therein). The pulsation is found to be in phase with that of the middle-band X-ray emission (eg. 2.5--5.0 keV band; see Figure~3). However, in case of Her~X-1, the pulsating soft component is found to be shifted by about 230$^\circ$ in phase from the power-law continuum component (Endo et al. 2000). A systematic and detailed study of the accreting X-ray pulsars which show soft excess revealed that the soft excess is a common feature in the X-ray pulsars (Hickox et al. 2004). Both the pulsating and non-pulsating natures of the soft excess seen in various X-ray pulsars suggest a different origin of emission of the soft components. The possible origins of the soft excess in accretion powered X-ray pulsars are (a) emission from accretion column, (b) emission by a collisionally energized cloud, (c) reprocessing by a diffuse cloud, and (d) reprocessing by optically thick material in the accretion disk (Hickox et al. 2004). The soft excess emission from the accretion column and by the reprocessing of harder X-rays by optically thick material in the accretion disk, are expected to show pulsations whereas in other cases, it is non-pulsating in nature. The pulsation of the soft component (blackbody component) in phase with the middle-band X-ray emission in A0535+262 suggests that the most probable region of soft X-ray emission is the accretion column and/or the inner accretion disk. The absorption column seems to have significant contribution to produce the dip like structure at soft X-rays at phases 0.65-0.80 in the count rate profile (left top panel of Figure~\ref{phrs}). Column density takes large values, $\sim\!1.2\times10^{22}$ atoms cm$^{-2}$, at phase 0.7-0.8. This may contribute to produce the dip like structure in the profiles below 1 keV which disappears from the profiles at higher energies. On the other hand, the primary dip exists at all the energies. The observed primary dip in the pulse profile can be understood by the change in the NPEX model parameters over pulse phase. The NPEX continuum model is an approximation of the unsaturated thermal Comptonization in hot plasma (Makishima et al. 1999). At low energies, it reduces to the ordinary power-law with negative slope. The lower value of the power-law index ($\alpha_1$) at the primary dip phase implies large optical depth to the Compton scattering. This means that many photons are scattered-out from the line of sight. It explains the presence of the primary dip in the soft and hard X-ray pulse profiles at same phase. In addition to the primary dip, the deepening of the dip at the phases 0.0-0.17 (Figure~\ref{erpp}) can also be understood by the changes of the parameters of the Comptonizing plasma. The power-law index ($\alpha_1$) tends to be larger and the exponential cut-off energy is smaller in this phase range. This means that the Comptonizing plasma has a relatively smaller optical depth and a lower temperature. This reduces the efficiency of Compton up-scattering and the number of hard photons, which causes the deepening of the dip at phases 0.0-0.17 above 5 keV. | 7 | 10 | 0710.0961 |
0710 | 0710.4545_arXiv.txt | We present novel evidence for a fine structure observed in the net-circular polarization (NCP) of a sunspot penumbra based on spectropolarimetric measurements utilizing the Zeeman sensitive \FeI\ 630.2\,nm line. For the first time we detect a filamentary organized fine structure of the NCP on spatial scales that are similar to the inhomogeneities found in the penumbral flow field. We also observe an additional property of the visible NCP, a zero-crossing of the NCP in the outer parts of the center-side penumbra, which has not been recognized before. In order to interprete the observations we solve the radiative transfer equations for polarized light in a model penumbra with embedded magnetic flux tubes. We demonstrate that the observed zero-crossing of the NCP can be explained by an increased magnetic field strength inside magnetic flux tubes in the outer penumbra combined with a decreased magnetic field strength in the background field. Our results strongly support the concept of the uncombed penumbra. | \label{sec:intro} The asymmetry of the Stokes parameters with wavelength is a powerful diagnostic tool to probe the relation between flows and the magnetic field in the atmosphere of sunspots. Line asymmetries convey information about lateral (within one resolution element) and the line-of-sight (LOS) gradients and discontinuities of physical parameters. Since the first systematic characterization and interpretation of these asymmetries and their spatial distribution throughout a sunspot by \citet[and references therein]{sanchezalmeida+lites1992} the complementary improvements achieved in both observational techniques and theoretical modeling, have advanced our understanding of particularly the penumbral inhomogeneities. It is general consent that a combination of non-trivial magnetic field and flow topologies along the LOS are the cause of the observed penumbral Stokes asymmetries. A simple and instructive way to characterize the asymmetry of Stokes V profiles is through the net-circular polarization (NCP), defined as $\mathcal{N}=\int V(\lambda) d\lambda$, where the integral is performed over one spectral line only. The sum of the NCP of individual lines than leads to the well-known broad-band circular polarization (BBCP) first noticed by \cite{illing+landmann+mickey1974a, illing+landmann+mickey1974b}. The appearance of the NCP is very distinct in the penumbra and shows a peculiar difference when between observations in the visible or in the near infrared wavelength region: in the visible at \FeI\,630.2\,nm the NCP maps tend to be symmetric about the disk center and sun center connection line \citep{sanchezalmeida+lites1992, westendorp+etal2001b, mueller+etal2002} while infrared NCP maps at \FeI\,1564.8\,nm reveal an antisymmetric distribution \citep{schlichenmaier+collados2002, bellot+balthasar+collados2004}. \cite{schlichenmaier+etal2002} and \cite{mueller+etal2002} demonstrate that the uncombed penumbral atmosphere \citep{solanki+montavon1993} mimicked by the \textsl{moving tube model} \citep{schlichenmaier+jahn+schmidt1998a} reproduces the azimuthal behavior of the NCP and, furthermore, convincingly show that the incisive discrepancy between the visible and the near infrared can be understood when the effects of anomalous dispersion \citep{landolfi+landi1996} are included. So far penumbral NCP maps of the \FeI\,630.2\,nm line have been only presented before by \cite{sanchezalmeida+lites1992} (in form of a contour line map), \cite{westendorp+etal2001b}, \cite{mueller+etal2002} and \cite{mueller+etal2006}. These maps visualize very well the symmetric behaviour of the NCP but lack the spatial resolution to reveal the presence of any potential fine structure. In this Letter we focus on high-spatial resolution spectropolarimetric measurements of a sunspot to investigate inhomogeneities in the penumbral magnetic field and flow field. Instead of utilizing inversion techniques we calculate the NCP and line-of-sight (LOS) velocities and elaborate on the properties of the same. In order to interpret the observational results we use the 3D geometric flux tube model VTUBE \citep{mueller+etal2006} and perform subsequent radiative transfer calculations from which NCP and Doppler maps can be synthesized. | \label{sec:conclusions} In this paper, we have presented novel evidence for a spatial fine structure in the visible NCP of the \FeI\,630.25\,nm line. We find for the first time that the NCP is structured in radial filaments very much alike maps of other line parameters, e.g. the intensity, equivalent width and line width \citep[see e.g.][]{johannesson1993, rimmele1995a, rimmele1995b, tritschler+etal2004}. The spatial scales of the NCP filaments are similar to those found in the observed LOS velocity map although we cannot find a good overall correlation between the signals of the NCP and the LOS velocity. We conjecture that higher spatial resolution is needed to settle this issue. \cite{martinezpillet1997} find a correlation between the azimuthal variation of the NCP and the magnetic field inclination on angular scales of $10^{\circ}$. Although we cannot provide information about the inclination at this point we find that our observations show azimuthal variations of the NCP that are only of the order of $4^{\circ}$. Furthermore, we have observed a zero-crossing of the NCP on the outer center-side penumbra and have demonstrated that this signature can be explained by an increased magnetic field strength inside magnetic flux tubes in the outer penumbra, while the magnetic field strength is decreased in the background field ($\Delta B$-effect). To this end, we have compared high-resolution observations with maps of the NCP and the LOS velocity that have been synthesized from a 3D geometric sunspot model that has been coupled with a 1D radiative transfer code \citep[see][ for details]{mueller+etal2006}. We point out that the space of possible atmospheric models that can reproduce the observations is significantly constrained by taking into account the information contained in Doppler maps in addition to the NCP maps. This paper extends the work by \cite{solanki+montavon1993}, \cite{martinezpillet2000}, \cite{mueller+etal2002}, and \cite{mueller+etal2006} and clearly demonstrates the need to model the "uncombed" penumbra of sunspots in greater detail. This would include also taking into account the effect of field line curvature around flux tubes as considered by \cite{borrero+bellot+mueller2007}. | 7 | 10 | 0710.4545 |
0710 | 0710.1010_arXiv.txt | The evolution of single stars at low metallicity has attracted a large interest, while the effect of metallicity on binary evolution remains still relatively unexplored. We study the effect of metallicity on the number of binary systems that undergo different cases of mass transfer. We find that binaries at low metallicity are more likely to start transferring mass after the onset of central helium burning, often referred to as case C mass transfer. In other words, the donor star in a metal poor binary is more likely to have formed a massive CO core before the onset of mass transfer. At solar metallicity the range of initial binary separations that result in case C evolution is very small for massive stars, because they do not expand much after the ignition of helium and because mass loss from the system by stellar winds causes the orbit to widen, preventing the primary star to fill its Roche lobe. This effect is likely to have important consequences for the metallicity dependence of the formation rate of various objects through binary evolution channels, such as long GRBs, double neutron stars and double white dwarfs. | During their life stars can expand to a radius which is up to 1000 times bigger than their initial radius. In close binary systems, they start to transfer mass if their radius exceeds a critical radius, the Roche lobe radius. It is usually the initially most massive star, the primary star, which ``fills its Roche lobe'' first. At solar metallicity binary evolution has been studied by various groups, while their evolution at low metallicity is relatively unexplored. Heavy elements, such as carbon, oxygen, nitrogen and iron, are important contributors to the opacity of stellar material in metal rich stars. Due to the lower opacity, metal poor stars are generally hotter and more compact. As it is the evolution of the radius of a star in a binary which determines if and when mass transfer occurs, we expect that the evolution of binaries in metal poor environments is significantly different from binaries in the solar neighborhood. In this work we study the radius evolution of stars with different masses in order to infer the frequency of various cases of binary evolution as a function of metallicity. In a second contribution to this proceedings \citep{acc07} we discuss the effect of metallicity on the expansion of accreting main sequence stars and the potential consequences for binary evolution. \begin{figure} \includegraphics[bb=190 0 900 500, clip, width=\columnwidth]{deminkea_poster2_fig1.eps} \caption{ Evolution track of a 6 \Msun star at solar (Z = 0.02) and low metallicity (Z =0.00001). \label{hrd} } \end{figure} \begin{figure*} \includegraphics[ bb=50 10 680 480, width=\columnwidth]{deminkea_poster2_fig2.ps} \includegraphics[ bb=50 10 680 480, width=\columnwidth]{deminkea_poster2_fig3.ps} \caption{ The initial orbital periods that lead to case A, B and C mass transfer are given as function of the primary mass assuming an initial mass ratio of 0.75 \citep[inspired by][]{Webbink79}. Widening of the orbit due to angular momentum loss in the form of stellar winds is taken into account. The hatching indicates approximately in which cases the donor has a convective envelope. \label{cases} } \end{figure*} | We find that at low metallicity the donor star in a binary is more likely to have developed a massive CO core before the onset of mass transfer. For binaries with primary masses above 30\Msun Case C mass transfer may occur almost exclusively at low metallicity. This effect potentially has important consequences for the metallicity dependence of the formation rate of various objects through binary evolution channels such as long Gamma Ray Burst progenitors \citep[see discussion section of][]{wolf+podsiadlowski07} but also double neutron stars and double white dwarfs. More work on binary evolution at low metallicity is needed. The precise fraction of Case C binaries as function of metallicity depends on assumptions about the distribution of initial binary parameters, the description of mixing and overshooting and the (metallicity dependence) of the mass loss rate. We find that the maximum radius can vary by more than a factor of 10 when we compare calculations of the STARS code to calculations done with the code described in \cite{Yoon+ea06} and when we vary the amount of overshooting. In a forth coming paper we will present the range case A, B and C mass transfer for a range of metallicities for different assumptions for the uncertain parameters \citep{CASEABC08}. \begin{theacknowledgments} We would like to thank Arend-Jan Poelarends, Matteo Cantiello, and Philip Podsiadlowski for interesting discussions and useful comments and suggestions. \end{theacknowledgments} | 7 | 10 | 0710.1010 |
0710 | 0710.3223_arXiv.txt | Recent advances have made it possible to obtain two-dimensional line-of-sight magnetic field maps of the solar corona from spectropolarimetric observations of the \ion{Fe}{13} 1075 nm forbidden coronal emission line. Together with the linear polarization measurements that map the azimuthal direction of the coronal magnetic field projected in the plane of the sky containing Sun center, these coronal vector magnetograms allow for direct and quantitative observational testing of theoretical coronal magnetic field models. This paper presents a study testing the validity of potential-field coronal magnetic field models. We constructed a theoretical coronal magnetic field model of active region AR 10582 observed by the SOLARC coronagraph in 2004 by using a global potential field extrapolation of the synoptic map of Carrington Rotation 2014. Synthesized linear and circular polarization maps from thin layers of the coronal magnetic field model above the active region along the line of sight are compared with the observed maps. We found that the observed linear and circular polarization signals are consistent with the synthesized ones from layers located just above the sunspot of AR 10582 near the plane of the sky containing the Sun center. | \label{sec:intro} Understanding the static and dynamic properties of the solar corona is one of the great challenges of modern solar physics. Magnetic fields are believed to play a dominant role in shaping the solar corona. Current theories also attribute reorganization of the coronal magnetic field and the release of magnetic energy in the process as the primary mechanism that drives energetic solar events. However, direct measurement of the coronal magnetic field is a very difficult observational problem. Early experiments have demonstrated the feasibility of the measurement of the orientation of the coronal magnetic fields by observation of the linear polarization of forbidden coronal emission lines (CELs) in the visible and at IR wavelengths \citep{eddy1967,mickey1973,arnaud1982, querfeld1982,tomczyk_et_al_2007}. Radio observations have also been successful in measuring the strength of the coronal magnetic field near the base of the solar corona \citep[e.g.,] [and references therein] {brosius_2006}. Direct measurement of the coronal magnetic field strength at a higher height by IR spectropolarimetry of the CELs was achieved only recently \citep{lin_et_al_2000, lkc_2004}. Without direct measurements, past studies involving coronal magnetic fields have relied on indirect modeling techniques to infer the coronal magnetic field configurations, including coronal intensity images observed in the EUV and X-ray wavelength ranges and numerical methods that reconstruct the three-dimensional coronal magnetic field structure by extrapolation and MHD (magnetohydrodynamics) simulation based on photospheric magnetic field measurements. Since experimental verification of theories and models is one of the cornerstones of modern science, the lack of observational verification of these indirect magnetic field inference methods that are in widespread use is a very unsatisfactory deficiency in our field. Our 2004 observations were obtained above active region AR 10582 right before its west limb transit. We have obtained the first measurement of the height dependence of the strength of the line-of-sight (LOS) component of the coronal magnetic field in this data, and it showed an intriguing reversal in the direction of the LOS magnetic field at a height of approximately $0.15\ R_{\odot}$ above the solar limb, as shown in Figures 4 and 5 of \cite{lkc_2004}. This feature and the observed linear polarization map have presented us with our first opportunity to carry out a comprehensive observational test of our coronal magnetic modeling methods in which the strength and direction of the magnetic fields predicted by the models can be directly checked by the observations. Force-free extrapolation of photospheric magnetic fields is currently the primary tool for the modeling of coronal magnetic fields. However, it is not without limitations or uncertainties. For example, the force-free assumption does not hold true in the photosphere and low chromosphere, and possibly in the high corona above 2 $R_{\odot}$ \citep{gary_2001}. Moreover, a different assumption (current-free, linear and nonlinear force free) about the state of the electric current in the corona can lead to substantially different extrapolation results. Without direct magnetic field measurements, many fundamental questions concerning the basic assumptions and the validity of our tools cannot be addressed directly. To date, questions like ``Is a potential field approximation generally an acceptable approximation for coronal magnetic fields?" or ``Do linear or nonlinear force-free extrapolations provide a more accurate description of the coronal magnetic field?" can only be addressed by visual comparison between the morphology of selected field lines of the extrapolated magnetic field model and observational tracers of coronal magnetic fields such as the loops seen in EUV images. However, these visual tests are qualitative and subjective. Furthermore, they assume the coalignment between the magnetic field lines and the loops in EUV images, which has not been verified observationally. As our first test, we attempted to address the question ``Is the potential field extrapolation generally an acceptable approximation for the coronal magnetic field?" This test was conducted by comparing the observed polarization maps of AR 10582 to those derived from a coronal magnetic field model constructed from the potential field extrapolation of photospheric magnetic field data. However, before we present our study, we should point out that because of the nature of our modeling tools and observational data, the results and conclusions of this research are subject to certain limitations and uncertainties. One of the intrinsic limitations of the observational data used in this study is that because of our single sight line from Earth to the Sun, the photospheric and coronal magnetic field observations cannot be obtained simultaneously. In the case of global coronal magnetic field models, the whole-Sun photospheric magnetic field data used as the boundary condition of the extrapolation can be obtained only over an extended period of time. Although the large-scale magnetic structure of nonflaring active regions may appear stable over a long period of time, high-resolution EUV observations have shown that the small-scale coronal structures are constantly changing. Thus, studies such as ours that compare coronal magnetic field observations and models constructed from photospheric magnetic fields inevitably are subject to uncertainties due to the evolution of the small-scale structures in the regions, and we should not expect a precise match between the observed and synthesized polarization maps. Another deficiency due to the single sight line of our observations is the lack of knowledge of the source regions of the coronal radiation. This is perhaps the most limiting deficiency of the observations and models of this research. The uncertainty of the location of the source regions associated with coronal intensity observations due to the low optical density of the coronal plasma and the resulting long integration path length is familiar. In the case of the coronal magnetic observations, the LOS integration problem prevents us from performing an inversion of the polarization data to reconstruct the three-dimensional magnetic field structure of the corona for direct comparison with those derived from extrapolations or MHD simulations. On the other hand, since extrapolation techniques do not include the thermodynamic properties of the plasma in the construction of the coronal magnetic field models, they do not include information about the location of the CEL source regions either. Therefore, they cannot predict the intensity and polarization distribution of the CEL projected on the plane of the sky that are needed for direct comparison with the polarimetric observations. Without the information about the location of the source regions from the observational data and the extrapolated models, we adopted a trial-and-error approach in which synthesized linear and circular polarization maps were derived using empirical source functions and the extrapolated potential magnetic field model, and were compared directly with those obtained from observations. Obviously, if acceptable agreement can be achieved with any of the models tested, then we can argue with a certain degree of confidence that these models are plausible models of the observed corona and that the potential field approximation is a reasonable approximation of the coronal magnetic fields. Nevertheless, we should emphasize that this trial-and-error approach is not an exhaustive search of all the possible source functions and therefore cannot provide a clear-cut true-or-false answer. In other words, even if no acceptable agreement can be found with all the model source functions we have considered, the potential field approximation still cannot be dismissed completely. | This research examines observationally the validity of current-free, force-free potential-field approximation for coronal magnetic fields. We conducted a study comparing observed and synthesized spatial variations of the linear and circular polarization maps in the corona above active region AR 10582, which after a week of extensive flaring activities should have settled into a minimum energy configuration that could be adequately modeled by a potential field model. The coronal magnetic field model used for this study was constructed from a global potential field extrapolation of the synoptic photospheric magnetogram of Carrington cycle 2014 obtained by the \SOHO/MDI instrument. Because the most important source of error of this type of study is the uncertainty of the location of the source function of the coronal radiation, we first tested three analytical, but empirically determined, source functions. These simple source functions are based on a gravitationally stratified atmospheric density model with a uniform temperature in the entire modeled volume, supplemented by magnetic weighting functions based on the observational impression that, at least at the length scale of the typical active region size, CEL radiation seems to be correlated with the strength of the photospheric magnetic fields. We found that none of these empirical source functions can adequately reproduce the observed linear polarization maps, although it seems that the source functions that include both density and magnetic fields produced slightly better results. Based again on the observational impression that the coronal intensity structures have a spatial scale much smaller than that of the active regions, we then compared the observed polarization maps with those constructed from thin (56 Mm FWHM) layers along the LOS. In this analysis, we found that polarization maps originating from layers located near the sunspot of the region are in reasonable agreement with the observed ones. However, the best fit for linear and circular polarization did not occur at the same layer. They are separated by a distance of about 50 Mm. Does the small discrepancy between the best-fit locations of the linear and circular polarization weaken the support for the potential field extrapolation? As we have discussed in \S\ref{sec:intro}, many uncertainties conspire to limit the precision of this study. For example, the difference may be due to the evolution of the small-scale photospheric magnetic field of the active region, and we do not have any observational information that we can use to test this possibility. The assumption of uniform temperature distribution in our source function is certainly not a physically realistic assumption. Therefore, it is not possible to assess the significance of the small difference in the location of the source regions of the linear and circular polarization. Furthermore, in addition to density and temperature, the source functions of CEL linear and circular polarization depend on different components of the coronal magnetic field. So it is in fact physically reasonable that we would find the linear and circular polarization signals originate from slightly different locations. Finally, because the three parameters ($\sigma_p$, $\sigma_\chi$, and $H_0$) we used to evaluate the quality of the fit were obtained independently, the statistical significance that all three parameters reached minimum at approximately the same location near the strongest photospheric magnetic feature of the active region cannot be dismissed as pure coincidence. These considerations lead us to conclude that potential field extrapolation can be used to provide a zero-order approximation of the real solar corona if the active region is in a relatively simple and stable configuration. Additionally, this study suggests that, at least for isolated active regions, CEL radiation may originate from a region close to the strongest photospheric magnetic feature in the active region with a small spatial scale comparable to the characteristic size of the coronal loops seen in the intensity images. If this is confirmed, then a single-source inversion to infer the magnetic field directly from the polarimetric observation such as that proposed by \cite{judge_2007} may be justified. Our conclusion about the viability of potential field extrapolation as a coronal magnetic field modeling tool for stable active regions is supported by a study by \cite{riley_et_al_2006}, in which the coronal magnetic field configuration derived from a potential field model was found to closely match that derived from a MHD simulation in the case of untwisted fields. Nevertheless, we should emphasize that our conclusion is derived from a single observation of a simple and stable active region. Clearly, more observations and model comparison are needed for a more comprehensive test of this result. Can the potential field approximation be used to model more complicated active regions? Using radio observations, \cite{lee_et_al_1999} found that a force-free-field model yields better agreement between the temperatures of two isogauss surfaces connected by the modeled field lines of an active region with strong magnetic shear. This study thus provides observational evidence against the use of potential field approximations for the modeling of complex active regions. Therefore, linear and nonlinear force-free extrapolations should be employed in future testing of theoretical coronal magnetic field models using the IR spectropolarimetric observations to study if these models can offer a better description of the observed coronal fields. Can we distinguish the potential coronal magnetic field configurations from the non-potential ones with the spectropolarimetric observations of the coronal emission lines? In a numerical study, \cite{judge_2007} has demonstrated the sensitivity of LOS-integrated coronal polarization measurements to the electric current in the corona using theoretical coronal magnetic field models as input. Therefore, we should expect to find better agreement between observed and synthesized polarization maps for more complex active regions with linear or nonlinear force-free magnetic field models. Work to model AR10582 using the force-free extrapolation method is already underway, and we should be able to address this question in the near future. Since all extrapolation methods are subject to the ambiguities problem of the source regions, we will also employ MHD simulations that include both the magnetic and thermodynamic properties of the corona in the calculation to help resolve this problem. These are research activities that we will be pursuing in the near future as the solar cycle evolves toward the next solar maximum and more coronal magnetic field data become available. The greatest difficulty of this study is the uncertainty of the location of the source function due to the long integration path along the LOS. However, this is not a difficulty affecting only the interpretation of coronal magnetic field measurements. It affects the intensity observation as well, and is the primary reason that years into the operation of \SOHO/EIT and \TRACE, we still cannot deduce 3-D intensity and temperature structure of the corona using data from these instruments. Fortunately, this deficiency in our observing capability may finally be removed with the recent launch of the \STEREO\ mission (Solar TErrestrial RElations Observatory). For the resolution of the LOS integration problem in polarimetric observations, \cite{kramar_et_al_2006} have demonstrated the promising potential of vector tomography techniques. While stereoscopic coronal magnetic field observations will not be realized any time soon, this method can be applied to observations obtained over periods of several days during the limb transit of active regions, provided that the active regions are in a stable condition. This is perhaps the best observational tool available for the resolution of the LOS integration problem in the near future. | 7 | 10 | 0710.3223 |
0710 | 0710.5225_arXiv.txt | {Maser emission from the \hzo\ molecule probes the warm, inner circumstellar envelopes of oxygen-rich red giant and supergiant stars. Multi-maser transition studies can be used to put constraints on the density and temperature of the emission regions.} {A number of known \hzo\ maser lines were observed toward the long period variables R Leo and W Hya and the red supergiant VY CMa. A search for a new, not yet detected line near 475 GHz was conducted toward these stars.} {The Atacama Pathfinder Experiment telescope was used for a multi-transition observational study of submillimeter \hzo\ lines.} {The $5_{33} - 4_{40}$ transition near 475 GHz was clearly detected toward VY CMa and W Hya. Many other \hzo\ lines were detected toward all three target stars. Relative line intensity ratios and velocity widths were found to vary significantly from star to star. } {Maser action is observed in all but one line for which it was theoretically predicted. In contrast, one of the strongest maser lines, in R Leo by far \textit{the} strongest, the 437 GHz $7_{53}-6_{60}$ transition, is not predicted to be inverted. Some other qualitative predictions of the model calculations are at variance with our observations. Plausible reasons for this are discussed. Based on our findings for W Hya and VY CMa, we find evidence that the \hzo\ masers in the AGB star W Hya arise from the regular circumstellar outflow, while shock excitation in a high velocity flow seems to be required to excite masers far from the red supergiant VY CMa.} | Discussion} \subsection{\label{excitation}Modeling the excitation of water maser lines} Earlier excitation studies aimed at explaining stellar (and interstellar) \hzo\ maser pumping only considered a limited number of energy levels, both, due to limitations of computational capabilities and because collision rates were not known for excitation to and from higher energy states \citep[see, e.g., ][]{deJong1973, Deguchi1977, CookeElitzur1985}. Nevertheless, all these studies predicted strong inversion of the 22.2 GHz line, which was the only \hzo\ maser line known at the time. These studies also predicted maser action in a number of other transitions that were later indeed found to be masing, despite the fact that in some cases the collisional excitation rate coefficients used were vastly different from the ``modern'' values computed by \citet{Palma_etal1988}. The more recent, comprehensive, study by NM91 considered all 349 states of ortho and para water up to an energy of 7700 K above the ground state and extrapolated the Palma et al. rates to high enough temperatures. \citet{NeufeldMelnick1990} and NM91 performed statistical equilibrium calculations modeling vibrational ground state \hzo\ excitation and radiative transfer for plane parallel media under the large velocity gradient (LVG) assumption. \citet{NeufeldMelnick1990} specifically explore conditions under which, both, the the 22.2 GHz and the 321 GHz lines are masing for the case of an O-rich evolved star with a mass-loss rate ($3~10^{-5}$ \Mspy) intermediate between the values of our objects. They find that the pumping for these lines is dominated by collisions and that radiation plays a minor role. The same was also found for other transitions in circumstellar envelopes \citep[][ NM91]{Deguchi1977, CookeElitzur1985}. NM91 explored larger regions of (essentially) density and temperature parameter space in which the energy levels of certain transitions become inverted. They calculate maser emissivities and optical depths as functions of temperature and a variable, $\zeta'$, which is equal to a geometric factor, $G$, divided by the velocity gradient, $d{\rm v}/dr$, times the product of the hydrogen nuclei, $n$, and the water density, $n({\rm H}_2{\rm O})$, i.e., $n\times n({\rm H}_2{\rm O})$ $ = n^2 x({\rm H}_2{\rm O})$, where $x({\rm H}_2{\rm O})$ is the \hzo\ abundance. $\zeta'$ is, thus, given by: $$\zeta' = G{n^2 x({{\rm H}_2{\rm O})}\over{d{\rm v}/dr}}$$ Assuming a plausible, constant, value for the \hzo\ abundance and the velocity gradient of $10^{-4}$ and $10^{-8}$~cm~s$^{-1}$~cm$^{-1}$, respectively, makes $\zeta'$ dependent on the H density squared alone. Measuring the H density in units of $10^9$ \ccm\ in the parametrization of NM91, values of $log_{10}\zeta' = -1, 0,$ and 1 correspond to H densities of 0.32, 1, and $3.2\times10^9$ \ccm. \subsection{\label{comparison}Comparison with the APEX observations} NM91 predicted maser emission in all of the lines listed in Table \ref{lines}, except for the 437 GHz $7_{53}-6_{60}$ and 443 GHz $7_{52} - 6_{61}$ transitions in, roughly, this range of values quoted above. Maser emission in all their predicted maser lines has indeed been observed in star-forming regions and/or red giant stars, except for the 355 GHz $17_{4,13} - 16_{7,10}$ transition. They present opacities and emissivities (which are proportional to isotropic photon luminosities) for assorted lines calculated for two values of the kinetic temperature, 400 K and 1000 K. In particular, they find the very widespread 22.2 GHz \kbw\ transition to be by far the most luminous and the most robust maser line in the sense that it is inverted over the widest range of physical conditions in particular at extreme values of the density. A general trend is that for 1000 K the maximum emissivities in all the lines are reached for a factor of a few higher densities than at 400 K. In the calculations of NM91, the peak emissivities of all lines but the 22.2 GHz line are similar within a factor of a few (although they occur at different values of $\zeta'$). At, both, 400 and 1000 K, the peak emissivity of the 22.2 GHz line is much higher than that of any of the others, by a factor of $\sim10$ at 400 K and $\sim5$ at 1000 K. The peak emissivity is also reached at an order of magnitude higher value of $\zeta'$ than for other lines. As can be seen from Table \ref{lines} and Fig. \ref{lineratios} in none of our stars does the 22.2 GHz line show the behavior predicted by these calculations, i.e. that it should be much more luminous than all the other lines. While it is among the strongest lines in VY CMa, it is anomalously weak in R Leo. We note that R Leo showed an even more extreme 321/22.2 GHz line ratio in 1989/1990 \citep{MentenMelnick1991}. In this star, the 22 GHz line is known to undergo extreme and erratic variations changing from not or barely detectable to hundreds of Jy within dozens of days \citep{Rudnitskij1987}. \subsection{The ''new'' 475 GHz transition} For the newly discovered 475 GHz $5_{33} - 4_{40}$ transition NM91 predict, as for the 439 and 471 GHz lines, maser action only for their 1000 K calculation. The presence of maser action in this line is, thus, an important indicator of a high temperature environment. \subsection{\label{predictions}Non- and mispredictions} NM91\textit{do not} predict maser emission for the 437 GHz $7_{53}-6_{60}$ transition, which we find to be by far the strongest line in R Leo and the second strongest in VY CMa. This line has only been found toward evolved stars and not toward SFRs \citep{Melnick_etal1993}. Pumping by the strong IR field in these objects, possibly via line overlaps, not considered by NM91, may be responsible for its excitation. NM91 \textit{do}, however, predict maser action in the extremely high excitation (5764.3 K) $17_{4,13} - 16_{7,10}$ transition near 355 GHz high with peak emissivities at $\zeta'$ values 1--2 order of magnitude higher than values for which other lines show their peak emissivities. For constant water abundance and velocity gradient this translates into H$_2$ densities of order $10^{10-11}$ cm$^{-3}$, which is 1--2 orders of magnitudes higher than the values over which most of the other lines show maser action. Our non-detection of this line (see Table \ref{lineresults} and \S\ref{comparison}) seems to indicates that under such extreme conditions the necessary requirements for producing detectable maser emission may not be met. \subsection{A note of caution} Guided by the described model calculations, observations of combinations of maser lines in principle should put constraints on the physical parameters of the masing regions. One naive hope might be: The more lines one finds masing the tighter the physical parameters of the masing region can be constrained. One might even hope that the observed relative luminosities of the lines might provide yet tighter constraints. Unfortunately at present, in the absence of more realistic models of the maser regions such expectations are most likely overly optimistic for several reasons: Comparing the emissivities or their ratios predicted by the generic calculations of NM91 for a \textit{plane-parallel(!)} medium with measured data probably does not make much sense. Moreover, an inherent assumption in the values calculated by NM91 is that all the lines are saturated throughout the region in which they are inverted. There is good observational evidence that this is not true for the 22.2 GHz line around W Hya \citep{ReidMenten1990} although this issue is unclear (and difficult to address) for other transitions and other stars. Another caveat when using line ratios as diagnostics is time variability. While all our APEX and Effelsberg data were taken within 1--2 weeks, at least the 22.2 GHz line in some sources (including R Leo) is known to show week-to-week variability (see \S\ref{comparison}). Information on variability on such small time scales is not available for the submillimeter lines, but presumably all of them are variable. This has been directly shown from a comparison of line profiles taken at different epochs 1--2 months apart for the 321 and 325 (and 22.2) GHz lines \citep{MentenMelnick1991,YatesCohen1996} Furthermore, to use multi-line data as diagnostics, an inherent assumption is that all lines arise from the same region (whose physical parameters are to be constrained). However, for VY CMa, the different shapes of the lines and the different velocity ranges they cover suggest that this is not true and that at least a significant portion of the lines' emission arises from disjoint regions. Clearly, high resolution interferometric observations are necessary to get a clearer picture of this. Given the high expected brightness temperatures, such observations with the Atacama Large Millimeter Array (ALMA) will afford studies of R Leo and W Hya with a resolution better than a stellar radius and a few stellar radii for VY CMa \citep[see \S4.4 of ][]{Menten2000}. All the lines discussed here have frequencies within ALMA's 275--370 and 385--500 GHz ``First Light'' Receiver bands \citep[``Bands 7 and 8'', ][]{Wilson_etal2005}. Observations of the 321 and 325 GHz lines can already be made today with the Submillimeter Array \citep{Ho_etal2004}. It is clear that current models for \hzo\ maser excitation in circumstellar envelopes fail to explain the observational picture. Velocity-resolved observations of a large number of maser and non-maser lines from a wide range of energies above the ground state with the Heterodyne Instrument for the Far Infrared (HIFI) aboard the Herschel satellite to be launched in the near future will soon deliver very tight constraints on \hzo\ excitation. Together with high spatial resolution data for high excitation (mostly maser) lines (see above), which will \textit{only} be attainable with ground-based interferometers and more sophisticated radiative transfer modeling these observations will lead to an understanding of water, which dominates the thermal balance in these regions. \subsection{\label{differentnature}Different natures of AGB star and red supergiant water masers?} Let us come back to the fact, mentioned in \S\ref{spectralappearance}, that toward W Hya, the 22.2 GHz \hzo\ maser emission arises from a ring with a diameter of 24 AU centered on the star. In contrast, toward VY CMa, the maser emission arises from a $770\times440$ AU region and proper motion observations appear to indicate that the masers partake in a bipolar outflow \citep{RichardsCohen1998}, although we notice that this flow would be directed perpendicular to the larger-scale CO outflow imaged by \citet{Muller_etal2007}. Are these distributions consistent with the masers arising from a ''normal'', i.e. steadily expanding circumstellar outflow? In the following, we investigate whether over the whole of the maser distributions in both objects the temperatures (calculated from the luminosities of the stars) and the densities (derived from the mass-loss rates) are conducive for maser excitation. For W Hya \citet{Justtanont_etal2005} modeled CO data to derive a temperature, $T$, profile, i.e., $T$ vs. distance from the star, $r$. For the maser region ($r = 1.8~10^{14}$ cm), one finds $T \approx 1000$ K from their Fig. 3. Also, assuming, like these authors, a luminosity of 5400 \Lsun\ and an effective temperature of 2500 K, respectively, one calculates a comparable 970 K at this $r$ from the Stefan-Boltzmann law. The H$_2$ density in a spherically symmetric outflow at distance $r$ from the central star is given by $$n({\rm H}_2) = 1.13~10^{43}\Mdot(\Mspy)r^{-2}({\rm cm}){\rm v}_{\rm exp}^{-1}(\kms)$$ where \Mdot\ and ${\rm v}_{\rm exp}$ are the mass-loss rate and the expansion velocity, respectively, \citep[see, e.g. ][]{Millar1988,MentenAlcolea1995}. If we assume for the masing region $n({\rm H}_2) = 10^9$ \ccm\ (see \S\ref{comparison}), $r = 1.6~10^{14}$ cm and an expansion velocity of 2 \kms\ and invert this equation to calculate the mass-loss rate, we obtain $\Mdot = 6~10^{-6}$ \Mspy, which is half the rate derived by \citet{ZubkoElitzur2000} and in between the estimates of \citet{Neufeld_etal1996} and \cite{Justtanont_etal2005} and, thus, hopefully, a plausible value. \footnote{\citet{Neufeld_etal1996} derive a large value of $0.3$--$2~10^{-5}$ \Mspy\ for \Mdot\ from modeling emission of thermally excited \hzo\ lines observed by the Infrared Space Observatory (ISO). In contrast, \citet{Justtanont_etal2005}, also using data from ISO as well as the Odin satellite, derive a much smaller mass loss rate of $(2.5\pm0.5)~10^{-7}$ \Mspy, comparable to values derived from CO lines. Their low \Mdot\ comes ``at the expense'' of an extremely high \hzo\ abundance of [\hzo/H$_2$ $=2~10^{3}$, which is significantly higher than the cosmic abundance of O. To explain it they invoke an influx of water from evaporating icy bodies in an Oort cloud analog, a mechanism that had earlier been advocated to explain the presence of water in IRC10216's envelope \citep{Melnick_etal2001}. Although the calculations of Justtanont et al. make a strong case in favor of a low \Mdot\ one wonders whether a higher \Mdot\ and a lower \hzo\ abundance in the range of the Neufeld et al. values might also be consistent with the measurements.} Doing the same exercise for the RSG VY CMa with a luminosity of $\approx 2~10^5$ \Lsun, we calculate for a region with an average radius of 300 AU (see \S\ref{spectralappearance}) a temperature of 480 K, which would be adequate at least for 22.2 GHz maser production. However, if we calculate the mass-loss rate required for a spherically symmetric outflow to maintain a density of $10^9$ \ccm\ at 300 AU, we obtain a value of 0.045 \Mspy. This is $\approx 200$ times higher than the generally quoted (already extremely high) mass-loss rate for this object (see \S\ref{intro}). We conclude that, in contrast to W Hya, at least the 22.2 GHz \hzo\ masers far from VY CMa cannot arise from the general material in a spherically symmetric outflow, but rather from special regions (e.g., shock fronts) within such a flow or from a collimated flow. This situation is reminiscent of \hzo\ masers in star-forming regions modeled by \citet{Elitzur_etal1989}. These authors invoke a temperature (400 K) similar to the value we derive for VY CMa. At this temperature, the highest excitation lines' maser emissivities are small (NM91), restricting their emission to regions closer to the star. This is in fact consistent with the smaller velocity ranges covered by the 321 and the 437 GHz lines. | 7 | 10 | 0710.5225 |
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0710 | 0710.2633_arXiv.txt | The near Main Sequence B stars show a sharp drop-off in their X-ray to bolometric luminosity ratio in going from B1 to later spectral types. Here we focus attention on the subset of these stars which are also Oe/Be stars, to test the concept that the disks of these stars form by magnetic channeling of wind material toward the equator. Calculations are made of the X-rays expected from the Magnetically Torqued Disk (MTD) model for Be stars discussed by Cassinelli \etal (2002), by Maheswaran (2003), and by Brown \etal (2004). In this model, the wind outflow from Be stars is channeled and torqued by a magnetic field such that the flows from the upper and lower hemispheres of the star collide as they approach the equatorial zone. X-rays are produced by the material that enters the shocks above and below the disk region and radiatively cools and compresses in moving toward the MTD central plane. It differs from the Babel \& Montemerle (1997) model in having a weaker B field and a large centrifugal effect. The dominant parameters in the model are the $\beta$ value of the velocity law, the rotation rate of the star, $S_o$, and the ratio of the magnetic field energy density to the disk gravitational energy density, $\gamma$. The model predictions are compared with the $ROSAT$ observations obtained for an O9.5 star $\zeta$ Oph from \Berghofer\ \etal (1996) and for 7 Be stars from Cohen \etal (1997). Two types of fitting models were used to compare predictions with observations of X-ray luminosities versus spectral types. In the first model we choose an estimate of the spin rate parameter $S_o$ from observed $\vsini$ values, and choose $\gamma$ using the threshold magnetic field value derived in Cassinelli \etal (2002), then the $\beta$ value is adjusted to fit the observations. For all but one case, the $\beta$ value was found somewhat larger than unity, a typical value derived for radially streaming stellar winds. This value, appropriate to a slowly accelerating wind, might be an indication that the magnetic field modifies the dynamics of the outflow from the star. In the second fitting model, we choose $\beta$ to be unity, $\gamma$ as in the first model, and adjust $S_o$ to fit the observations. For these comparisons with the X-ray observations we find that $S_o$ is in the range 0.49 to 0.88, which agrees with traditional estimates of the rotation rate of Be stars, but is below the breakup values that are required in recent non-magnetic models for Be star disks (Townsend \etal 2004). Extra considerations are also given here to the well studied Oe star $\zeta$ Oph for which we have $Chandra$ observations of the X-ray line profiles of the triad of He-like lines from the ion Mg XI. We find that a reasonably good fit is made to the observed Mg XI line profiles. The lines are predicted to form primarily at radial distance of about two stellar radii in the disk, and the ratio of the forbidden to intercombination lines (i.e., $f/i$) that is a diagnostic of source distances, agrees with this prediction. In addition, the lines are broad, with HWHM of about 400 $\kmsec$. Again this is in compatibility with the model predictions for the disk rotation of this star. Thus the X-ray properties add to the list of observables which can be explained within the context of the MTD concept. This list already includes the \Halpha\ equivalent widths and white light polarization of Be stars. We do not include here the disk density correction due to gravity, neglected in Cassinelli \etal (2002) but included in Maheswaran (2003). This will likely increase the \Halpha\ and polarization predictions but have little or no effect on the X-rays which are generated in the upstream region. Thus the X-rays are not sensitive to the cooled disk region where Keplerization of the disk material might be occurring. Nonetheless the process by which matter and angular momentum are added to the disk are equally important and this study indicates that X-ray properties are consistent with the overall MTD concept. | In a survey of the X-ray emission of near Main-Sequence B stars (or B V stars), Cassinelli \etal (1994) and Cohen \etal (1997) found a departure from the canonical ``law" relating X-ray luminosity to the bolometric luminosity for hot star X-rays: $L_x/L_B=10^{-7}$. This relation holds for stars throughout the O spectral range and extends to about B1 V. However, beyond that there is a sharp drop in the ratio values in going to B3 V by about 2 orders of magnitude. Cohen \etal (1997) investigated if the sharp decrease could be explained merely by the reduction of the wind outflow from these stars, and found that this could indeed explain the initial decrease in the X-ray luminosity, but that in going to even later B V stars another problem arose. The emission measures of the X-ray producing material at spectral type B3 V and later become larger than the predicted wind emission measures for B stars for the smooth wind case. Cohen \etal (1997) suggested that this could be the first indication that the late B stars lie at the transition to the outer atmospheric structure of cool stars, for which surface magnetic fields control the X-ray properties. These ideas were based on a spherical radial outflow picture for B stars. In fact, B stars are known to be rather rapid rotators. Bjorkman \& Cassinelli (1993) developed the Wind Compressed Disk (WCD) model, which suggested the wind from a rapidly rotating star would orbit towards the equatorial region where it would shock and compress the incident gas. The model had success in explaining the polarization properties of emission line Be stars (Wood \etal 1997). WCD was also supported, initially, by hydrodynamic simulations performed by Owocki, Cranmer \& Blondin (1994). However, in a more detailed consideration of the flow to the equator idea, Owocki, Cranmer \& Gayley (1996) found that non-radial line forces in a rotating and distorted star tend to impede the flow to the equator and produce a bipolar flow instead. Hence, the cause of the {\it disks} that exist around Be stars as opposed to {\it polar plumes}, has become a topic of much debate among theorists. Observers also found problems with the WCD idea. Hanuschik \etal (1996) found in their observations of {\it equator-on} Be stars, that the mass outflow speeds {\it detectable} in the disks were negligible compared with the steady but slow equatorial outflows predicted by the WCD model, though no one seems ever to have predicted whether in fact the WCD material is dense enough to be detectable in this way. Even more interesting was their observational conclusion that the azimuthal speed of the inner disk material was larger than the angular speed of the star from which the disk presumably originated (as it must be to remain in Keplerian orbit unless supported by other forces.) Specifically, observations of the equator-on Be star $\beta^1$ Mon showed that the Fe II line emission arising from the disk is broader than the $\vsini$ value derived from photospheric lines. Hanuschik \etal\ (1996) suggested that the equatorial disk material was in Keplerian motion about the star. In the context of any Keplerian paradigm, it is especially important to note that in order to form a Keplerian disk, there needs to be an increase in the specific angular momentum of the matter after it leaves the star. The mechanism for providing that additional angular momentum is currently an unresolved subject of debate (e.g. Baade \& Ud-Doula 2005, and Brown \& Cassinelli 2005). Transfer of mass and torquing of the outflow could be produced by magnetic fields rooted in the star's surface, as is well-known from magnetic rotator theory (Lamers \& Cassinelli 1999 Ch. 9). The existence of magnetic fields in hot stars is now well established (Donati \etal 2001, 2002). Thus, the Magnetically Torqued Disk (MTD) model was proposed by Cassinelli et al. (2002) for the disks around Be stars. In this model, the star is pictured as having a co-aligned dipolar field that both channels and torques the wind from the star towards a disk. Being that the detailed structure of the magnetic fields of these stars has not yet been determined, it seems reasonable to use the pure dipole as the most conservative hypothesis. Minimal magnetic fields were derived from the need to torque the outflow and the much denser disk material to Keplerian speeds, and the fields required were compared with upper and lower limits for hot star fields that had been derived by Maheswaran \& Cassinelli (1988, 1992). For stars of spectral class B2 V (which correspond to the most common class of Be stars), the field required to torque the dense disk is about 300 Gauss while that required to torque the wind is only around 10 Gauss. This difference is because the field needed to torque a wind is proportional to the square root of the density, and the density of the wind is about 3 orders of magnitude lower than the density at the equatorial plane in the disk. The fact that the higher figure is comparable to the fields that have now been derived from multi-line Zeeman effect measurements of the very slowly rotating star $\beta$ Cephei (Donati \etal 2001) shows that both the flow and the disk will in reality undergo MTD torquing for this star. In other stars, the magnetic fields (whose strengths have not yet been measured) may lie in an intermediate regime in which the wind material would be torqued to high specific angular momentum while being channeled, but have entirely Keplerian dynamics in the denser disk (Owocki \etal 2005). The MTD model was shown to produce the \Halpha\ emission observed in Be stars (Doazan \etal 1991), and also the level of intrinsic polarization seen (Quirrenbach \etal 1998). It is important to note that only a fraction of B stars are Be stars, and those stars identified as Be stars only spend a fraction of their time in a state with identifiable Be-star features. Therefore it is not necessary for a theoretical paradigm to cause a magnetically torqued disk for all possible sets of stellar and magnetic parameters, it is only necessary for a disk model to encompass a wide enough range of parameter space so as to make it reasonable that some stars show disks some of the time. Detailed comparison with observations will only be possible when the ``duty cycles'' of Be stars and the Be star fraction have been better determined observationally. In the original MTD model, the stellar wind mass flux and wind speed distribution were taken to be uniform over the stellar surface. However, in the case of a rapidly rotating star, this assumption is invalid since the rotation results in gravity darkening (Von Zeipel 1924) and reduces the wind mass flux and the terminal speed in the equatorial region (Owocki \etal 1998). By incorporating Gravity Darkening into the MTD model (MTDGD), Brown \etal (2004) derived several important disk properties such as, the dependence of the disk mass density distribution on its extent, the total number of disk particles, and the functional dependencies of the emission measure and polarization on the rotation rate ($S_o$) and wind velocity law ($\beta$). In contrast to what had been expected, they found that the critical rotation (or $S_o= 1$) is not optimal for creation of hot star disks. One important omission in the basic MTD formulation, as well as in MTDGD to date was the erroneous neglect of gravity in the disk density structure - see Brown \& Cassinelli (2005). Though $g_z \propto z$ is zero in the equatorial plane, its increase with $z$ enhances the density in the disk and causes it to grow with time. It will likely increase the \Halpha\ and polarization predictions. In turn, this will result eventually in radial escape of disk material, though whether this is slow and steady or episodic as claimed by Owocki \& ud-Doula (2003) is not yet clear. Other computational and analytic attempts to model disks similar to those envisioned here have met with mixed results. Owocki \& ud-Doula (2003) criticized the MTD idea because the model was found to be unstable in their MHD simulations. However, the numerical simulations by Keppens \& Goedbloed (1999, 2000), and Matt \etal (2000) demonstrated the existence of disks around some hot stars like post-AGB stars. Also Maheswaran (2003) has studied this scenario using analytic MHD and finds results that the magnetically torqued disks are likely to be persistent, which cast doubt on the numerical simulations of Owocki \& ud-Doula (2003). Thus further work is needed to test the basic ideas of the MTD and MTDGD models, either thorough numerical and analytic MHD calculations, or using observational diagnostics from radio to X-ray wavelengths. In this paper, however, we are primarily concerned with X-ray emission from the MTDGD models, emission which occurs well upstream of the dense disk, and explain the X-ray anomalies associated with Be stars. This will be little affected by the density in the disk itself though our use of MTD to find the outer disk radius will make our X-ray source emission measure estimates a little too high. A successful model should be able to explain the drop-off at B2 Ve, the apparently excessive X-rays of late BV stars, while using disk parameters consistent with theory and the entire set of observational data. This process can then be inverted to use X-ray properties of a star to derive limits on its wind, disk, and rotational properties. In Section 2, we describe how X-ray emission is produced by the model. The effects of model parameters on X-ray emission are discussed in Section 3. Comparisons of model predictions with both $ROSAT$ and $Chandra$ observations are presented in Section 4. The discussion and conclusions are presented in Section 5. | The magnetically torqued disk model including gravity darkening effects or MTDGD, has here been tested to see if it can explain the basic X-ray properties of Oe/Be stars for which the model was developed. We had already found in the original papers on the MTD concept by Cassinelli \etal (2002) and Brown \etal (2004) that the idea of mass in Be disks is channeled by magnetic fields is consistent with the \Halpha\ luminosities of Be stars and that the mass in the disks as derived from polarization observations is also explainable, though we again note that inclusion of $g_z$ will likely increase these in the model. We have used a model that explains how matter can enter a disk with sufficient angular momentum to explain the quasi-Keplerian disks of Be stars, and the model uses field strengths that are comparable to those being found for other B stars. However an essential property of the MTDGD model is that it requires that X-rays be produced owing to the abrupt braking of the wind at the shock fronts. In summary, the paper contains several interesting results. (a) The paper tests the prediction that X-rays should be produced by the impact of channeled winds onto a disk. These X-rays would probably not be predicted from other current Be star models such as those in which the disk is produced by an extraction of angular momentum from the surface of a critically rotating star. (b) The model was based on the assumption that Be stars are rotating at their traditional values of about 70 percent critical and these rotation rates were found to be sufficient to explain the X-ray emission within the context of the MTD picture. (c) The model was found to require fields of order $10^2$ Gauss for Be stars, and our results show that these are adequate for the broad band X-ray production. (d) Broad band $ROSAT$ X-ray fluxes can be produced from B1 to B8 with MTDGD model parameters. (e) Fits are achievable even for the late B8 stars without invoking the presence of a dwarf M companion. (f) The model predicted that the Helium-like ion Mg XI had an $\rm R_{fir}$ about 1.8 stellar radii, which agrees well with where that line emission is expected to arise in $\zeta$ Oph with $Chandra$ observations, i.e., the $fir$ lines are predicted to occur there and to be broad, with a half width of about 400 $\kmsec$ and the temperature structure is consistent with the formation of the Mg XI line. Our discussion thus far has dealt with understanding the fundamental properties of Be stars as revealed by the $ROSAT$ and $Chandra$ observations. We have raised several questions during this paper that can now be addressed. (1) For our latest star B7 IVe, we found that the observed $ROSAT$ level of X-rays could be produced if the velocity law had value $\beta =0.77$. This small value for $\beta$ means that the channeled wind is colliding with the disk at a larger fraction of terminal wind speed than is the case for our other stars, which have $\beta$ values ranging from 1.3 to 2.8. The other required parameters for this star seem plausible: $B \sim 50$ Gauss, $S_o \sim 0.6$, (in the first of our two fitting procedures). Cohen \etal (1997) suggested that the emission measure needed to explain X-rays from late B stars were excessive, and that perhaps X-rays from magnetically confined region at the base of the wind are needed. So from our model of this star it appears that a magnetically confined X-ray formation region at the base is not needed. A bipolar magnetic field could instead just be torquing and channeling the wind toward the disk via X-ray emitting shocks. (2) The sharp drop-off of the X-ray luminosity beyond about B2 V, also seems to be explainable with a plausible range of our $\beta$, $S_o$, and magnetic field values of about 685 to 130 Gauss for our B2 and B3 stars. It appears that the MTDGD model has the ability to answer two of the more difficult questions concerning existing X-ray observations of Be stars, with plausible parameters. The model can explain the X-ray luminosity across the B spectral band and it can explain reasonably well the observed line profile results from $Chandra$. Finally, it is important to note that our model results are not dependent on the nature of the high density regions on the equatorial plane for which there is controversy regarding field wrapping and magnetic breakouts. This is because our models are in effect providing information only about the X-ray formation regions at the boundaries of the disk, and the flow through these boundaries occurs before the matter reaches the cooled compressed region near the equatorial plane. We are not requiring that the gas be controlled by a strong field all the way to the equatorial plane, and in fact think that the gas is no longer dominated by the field in that region and is free to acquire a quasi-Keplerian orbital motion. The needed angular momentum had been transferred to the gas in the pre-shock magnetically torqued and channeled MTD flow. | 7 | 10 | 0710.2633 |
0710 | 0710.5822_arXiv.txt | Magnetar's persistent emission above 10 keV was recently discovered thanks to the imaging capabilities of the IBIS coded mask telescope on board the \int~ satellite. The only two sources that show some degree of long term variability are \zerosei~ and \rxs. We find some indications that variability of these hard tails could be the driver of the spectral variability measured in these sources below 10 keV. In addition we report for the first time the detection at 2.8 $\sigma$ level of pulsations in the hard X-ray tail of \zerosei. | Anomalous X-ray Pulsars (AXPs) and Soft Gamma-Ray Repeaters (SGRs) are two small classes of sources that are believed to be magnetar candidates, namely isolated neutron stars with magnetic fields larger than the quantum critical value B$_{QED}\sim$4.4$\times$10$^{13}$G. For a review of this class of objects see \cite{woodsrew}. They share some common properties like a long spin period in the range of $P$=2--12 s, a large period derivative of $\dot P$=10$^{-13}$--10$^{-10}$ s s$^{-1}$, and a typical X-ray luminosity of $L_{X}\sim 10^{34}$--10$^{36}$ erg s$^{-1}$. This luminosity is well above the rotational energy losses, and is believed to be powered by the decay of the huge magnetic field. Since the spectra below $\sim$10 keV are rather soft, the first \int~ detections above 20 keV of very hard high-energy tails associated with these objects came as a surprise \citep{kuiper,denhartog,rev,mere05,molkov,dg06}. The spectra flatten ($\Gamma\sim$ 1, where $\Gamma$ is the photon index) above 20 keV and the pulsed fraction of some of them reaches up to 100\% \citep{kuiper}. The discovery of these hard tails provides new constraints on the emission models for these objects since their luminosities might well be dominated by hard, rather than soft, X-rays. In this paper we will focus on the timing properties and long term variations of two magnetar candidates, \rxs~ and \zerosei, as measured by IBIS/ISGRI \cite{ibis,isgri} on board the \int~ satellite \cite{integral}. | 7 | 10 | 0710.5822 |
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0710 | 0710.2543_arXiv.txt | Transonic accretion flow with self-consistent treatment of the magnetic field is presented. My website \url{http://www.cfa.harvard.edu/~rshcherb/}. | The averaged quantities can be obtained in two different ways in magnetohydrodynamics. The first way is to solve 3D MHD equations and then average the results. The second way is to solve some system of equations on averages. Combination of numerical simulations and averaged theory brings phenomenology that can describe observations or experimental data. The problem of spherically symmetric accretion takes its origin from Bondi's work \citep{bondi}. He presented idealized hydrodynamic solution with accretion rate $\dot{M}_B.$ However, magnetic field $\vec{B}$ always exists in the real systems. Even small seed $\vec{B}$ amplifies in spherical infall and becomes dynamically important \citep{schwa}. Magnetic field inhibits accretion \citep{schwa}. None of many theories has reasonably calculated the magnetic field evolution and how it influences dynamics. These theories have some common pitfalls. First of all, the direction of magnetic field is usually defined. Secondly, the magnetic field strength is prescribed by thermal equipartition assumption. In third, dynamical effect of magnetic field is calculated with conventional magnetic energy and pressure. All these inaccuracies can be eliminated. In Section 2\ref{section_method} I develop a model that abandons equipartition prescription, calculates the magnetic field direction and strength and employs the correct equations of magnetized fluid dynamics. In Section 3\ref{results} I show this accretion pattern to be in qualitative agreement with Sgr A* spectrum models. I discuss my assumptions in Section 4 \ref{discussion}. | \label{discussion} The presented accretion study self-consistently treats turbulence in the averaged model. This model introduces many weak assumptions instead of few strong ones. I take dissipation rate to be that of collisional MHD simulations. But flow in question is rather in collisionless regime. Observations of collisionless flares in solar corona \citep{noglik} gives dissipation rate $20$ times smaller than in collisional simulations \citep{biskamp03}. However, flares in solar corona may represent a large-scale reconnection event rather than developed turbulence. It is unclear which dissipation rate is more realistic for accretion. Magnetic field presents another caveat. Magnetic field lines should close, or $\vec{\nabla}\cdot\vec{B}=0$ should hold. Radial field is much larger than perpendicular in the inner region. Therefore, characteristic radial scale of the flow is much larger than perpendicular. If radial turbulence scale is larger than radius, freezing-in condition does not hold anymore. Matter can freely slip along radial field lines into the black hole. If matter slips already at the sonic point, the accretion rate should be higher than calculated. Some other assumptions are more likely to be valid. Diffusion should be weak because of high Mach number that approaches unity at large radius. Magnetic helicity was found to play very small dynamical role. Only when the initial turbulence is highly helical, magnetic helicity conservation may lead to smaller accretion rate. Neglect of radiative cooling is justified a posteriori. Line cooling time is about $20$ times larger that inflow time from outer boundary. The study is the extension of basic theory, but realistic analytical models should include more physics. The work is underway. \begin{theacknowledgments} I thank my advisor Prof. Ramesh Narayan for fruitful discussions. \end{theacknowledgments} | 7 | 10 | 0710.2543 |
0710 | 0710.0546_arXiv.txt | Estimates of velocities from time series of photospheric and/or chromospheric vector magnetograms can be used to determine fluxes of magnetic energy (the Poynting flux) and helicity across the magnetogram layer, and to provide time-dependent boundary conditions for data-driven simulations of the solar atmosphere above this layer. Velocity components perpendicular to the magnetic field are necessary both to compute these transport rates and to derive model boundary conditions. Here, we discuss some possible approaches to estimating perpendicular flows from magnetograms. Since Doppler shifts contain contributions from flows parallel to the magnetic field, perpendicular velocities are not generally recoverable from Doppler shifts alone. The induction equation's vertical component relates evolution in $B_z$ to the perpendicular flow field, but has a finite null space, meaning some ``null'' flows, e.g., motions along contours of normal field, do not affect $B_z$. Consequently, additional information is required to accurately specify the perpendicular flow field. Tracking methods, which analyze $\partial_t B_z$ in a neighborhood, have a long heritage, but other approaches have recently been developed. In a recent paper, several such techniques were tested using synthetic magnetograms from MHD simulations. Here, we use the same test data to characterize: 1) the ability of the induction equation's normal component, by itself, to estimate flows; and 2) a tracking method's ability to recover flow components that are perpendicular to $\mathbf{B}$ and parallel to contours of $B_z$. This work has been supported by NASA Heliophysics Theory grant NNG05G144G. | \label{sec:intro} The large length scales and relatively high conductivity of the plasma in the solar corona imply that, to a good approximation, magnetic flux is frozen to the plasma there. Consequently, the coronal magnetic field is ``line-tied'' to the plasma in lower atmospheric layers where hydrodynamic forces can be stronger than Lorentz forces --- the photosphere and lower chromosphere --- and coronal evolution is strongly coupled to evolution in these layers. Accordingly, observations of magnetic field evolution below the Sun's corona --- typically, sequences of photospheric or chromospheric magnetograms --- provide crucial tools understand coronal evolution. Usually, vector magnetograms are more useful than line-of-sight (LOS) magnetograms for studying coronal evolution, because information derived from LOS measurements alone will not, in general, be consistent with the actual magnetic field, which has field components both parallel and transverse to the LOS. Although time series of vector magnetograms have historically been rare, SOLIS \cite{BW_Henney2002}, the Solar Optical Telescope (SOT, Tarbell 2006) \nocite{BW_Tarbell2006} on Hinode, and the Solar Dynamics Observatory's Helioseismic and Magnetic Imager (HMI) \cite{BW_Scherrer2005}, should dramatically improve photospheric vector magnetogram spatial and temporal coverage in the near future. Several techniques have been developed to derive flows from time series of magnetograms \cite{BW_Chae2001,BW_Kusano2002,BW_Welsch2004, BW_Longcope2004,BW_Georgoulis2006,BW_Schuck2006}. Estimated flows at the base of the corona can be used to derive the fluxes of magnetic helicity, energy, and free energy into the corona, \cite{BW_Demoulin2003,BW_Pariat2005,BW_Welsch2006}. Further, flow estimates can be used to provide time-dependent boundary conditions for data-driven simulations of coronal magnetic field evolution. This paper briefly reviews progress on estimating surface flows from magnetogram sequences, and demonstrates some aspects of the problem with test data. | \label{sec:discuss} We have briefly reviewed central concepts regarding surface flow estimation from magnetograms. In addition, we have presented the results of simple tests which demonstrate that the induction equation's normal component, equation (\ref{eqn:normal}), can be used to quantitatively estimate flux transport rates. Fluxes of magnetic energy and helicity derived from these estimates are, however, likely to possess significant systematic errors. We have also demonstrated that flow estimation techniques that depend upon evolution in $B_z$ alone are insensitive to flux transport along contours of $B_z$, compared to flux transport along $\nabla_h B_z$. Equations (\ref{eqn:dSdt}) and (\ref{eqn:dhdt_bf}), which depend on the dot products of $\mathbf{u} B_z$ with $\mathbf{B}_h$ and with $\mathbf{A}_P$ respectively, combined with the sensitivity of $\mathbf{u} B_z$ to $\nabla_h B_z$, suggest that if $\mathbf{B}_h$ and/or $\mathbf{A}_P$ lie primarily along $\nabla_h B_z$, then the fluxes of magnetic energy and/or helicity can, in principle, be recovered accurately from $\Delta B_z/\Delta t$. If, in contrast, $\mathbf{B}_h$ and/or $\mathbf{A}_P$ lie primarily along contours of $B_z$, then the fluxes of magnetic energy and/or helicity probably cannot be recovered accurately from evolution in $B_z$ alone. We note that, as discussed in \cite{BW_Welsch2007}, the ANMHD data used in the tests presented here differ from actual magnetograms in significant ways, so the properties of flows estimated from actual magnetograms will probably differ substantially from the properties of flows estimated from ANMHD data. | 7 | 10 | 0710.0546 |
0710 | 0710.2858_arXiv.txt | GRB 021206 is one of the brightest GRBs ever observed. Its prompt emission, as measured by RHESSI, shows an unexpected spectral feature. The spectrum has a peak energy of about 700~\keV\ and can be described by a Band function up to 4.5~\MeV. Above 4.5~\MeV, the spectrum hardens again, so that the Band function fails to fit the whole RHESSI energy range up to 17~\MeV. Nor does the sum of a blackbody function plus a power law, even though such a function can describe a spectral hardening. The cannonball model on the other hand predicts such a hardening, and we found that it fits the spectrum of GRB 021206 perfectly. We also analysed other strong GRBs observed by RHESSI, namely GRBs 020715, 021008, 030329, 030406, 030519B, 031027, 031111. We found that all their spectra can be fit by the cannonball model as well as by a Band function. | \label{sec:intro} The exact mechanism which produces $\gamma$-ray bursts (GRBs) has not yet been definitively established. Their prompt $\gamma$-ray spectra can be used to distinguish between different models. Several mathematical functions have been used for parametrizing the prompt $\gamma$-ray emission. Most commonly used is the empirical Band function \citep{Band93}, which is not motivated by a physical model. There have been attempts to distinguish between spectral models analysing the low energy part of the spectrum. \citet{Ghirlanda2003}, \citet{Ryde2004}, and more recently \citet{Ghirlanda2007} searched for blackbody components in GRB spectra with varying degrees of success. \citet{Preece2002}, using BATSE GRB spectra, tested the synchrotron shock model and conclude that it ''does not account for the observed spectra during the GRB phase''. Spectral studies above the peak energy are rare, one reason being the poor data quality because of lack of statistics. Combining BATSE and EGRET spectra, \citet{Nature2003} report a high energy component for GRB 941017. They find a photon index of about 1.0 at energies above 5~\MeV. In this paper we report a high energy component in GRB 021206 \citep{GCN_021206_IPN,GCN_021206_final}, observed with the Reuven Ramaty High Energy Solar Spectroscopic Imager RHESSI \citep{RHESSI}. Having a peak energy of about 700~\keV, the spectrum of this burst can be described by a Band function from 70~\keV\ up to 4.5~\MeV, with a high energy photon index $\beta \approx 3.2$. Above 4.5~\MeV, the spectrum hardens again, and can be described with a photon index $\beta' \approx 2.2$. This significant hardening around 4.5~\MeV\ can not be described with a Band function. But it seems to differ from the spectral hardening in GRB 941017 as well. There is one model that fits the entire RHESSI spectrum of GRB 021206: the cannonball model \citep{Dar2004,Dado2002,Dado2003}. The cannonball model predicts a spectral hardening at several times the peak energy with a high energy photon index reaching $\beta \approx 2.1$. The question immediately arises whether the cannonball model can improve our description of other GRB spectra. The difference between the Band function and the cannonball model arises only at the high energy part of the spectrum, where data usually suffer from low statistics. Therefore, we choose the strongest GRBs registered by RHESSI in the years 2002 to 2004. We find that they all can be fit by the cannonball model as well as by the Band function. The outline of the paper is the following: We first present shortly the instrument, the GRB selection, the spectrum extraction, and the fit method (\S \ref{sec:observations}). In the next section (\S \ref{sec:models}), many spectral functions are given. In \S \ref{sec:results}, the fit results for GRB 020715, GRB 021008, GRB 021206, GRB 030329, GRB 030406, GRB 030519B, GRB 031027, and GRB 031111 are presented. The fits are discussed and, if possible, compared to other measurements. The more general discussion, including an outlook, follows in \S \ref{sec:discussion}. We end with a short summary in \S \ref{sec:summary}. | \label{sec:discussion} \subsection{The spectral functions} What is an acceptable $\chi^2$? In the limit of many degrees of freedom ($n_{DoF} > 30$), $\chi^2$ is normal distributed with an expectation value of $n_{DoF}-0.5$ and a variance of $\sigma_{\chi^2} = \sqrt{2n_{DoF}-1}$. A fit is acceptable if $\chi^2$ is close to its expectation value {\em and} if the residuals scatter around zero over the whole fit range, i.e.\ if the fit ``looks good''. From Table~\ref{tab:chi2} we conclude that the \CB\ gives acceptable $\chi^2$ for {\em all} GRBs studied. And they also look good, as can be seen in Figs.~\ref{fig:020715} to \ref{fig:031111}. Except for GRB 021206, rear (Fig.~\ref{fig:spec_rear}), the same can be said for the Band function. In many cases, a cut off power law (CPL) fits the spectrum up to high energies, e.g.\ GRB 021008, GRB 030329 or GRB 031027. In these cases, Band and \CB\ improve the goodness-of-fit slightly, but all three spectral shapes fit the data acceptably. A broken power law fits sometimes, but usually not well. BBPL and BBmPL do not fit in general, BBPL worse than BBmPL. However, it should be mentioned that a blackbody component is expected---if at all---only at the beginning of a GRB (see e.g.\ \citet{Ryde2006} and references therein), whereas we fitted the entire duration of the bursts. When using BBPL, we often find that the PL component fits either at high energies or at low energies. This is also discussed by \citet{Ghirlanda2007} who studied six BATSE GRBs in detail, where low energy data from the WFC instrument (on board BeppoSAX) are available. They find that the WFC data fit the Band function or CPL extrapolation, but not the BBPL extrapolation to low energies. Arguing that the PL contribution is too simple, they try to fit a blackbody spectrum plus CPL. % We suggest to use our BBmPL function instead. Its modified PL component describes a spectral break from $dN/dE \propto E^{-(\beta-1)}$ at low energies to $dN/dE \propto E^{-\beta}$ at high energies. \subsection{\CB\ function } The present work is, to our knowledge, the first systematic attempt to fit the \CB\ function to prompt GRB spectra. The two terms in \refeq{eq:CB} have a simple meaning. According to the cannonball theory, all GRBs are associated with a supernova. The ambient light is Compton up-scattered by the cannonball's electrons, producing the prompt GRB emission. Some electrons are simply comoving with the cannonball, giving rise to the CPL term in \refeq{eq:CB}. Since the photon spectrum of the ambient light can be described by a thin thermal bremsstrahlung spectrum, $\alpha$ is expected to be $\approx 1$. The second term (mPL) is caused by a small fraction of electrons accelerated to a power law distribution, resulting in a photon index of $\beta \approx 2.1$. See e.g.\ \citet{Dado2005}, \S 3.8. or \citet{Dado2007}, \S 2 and 4.1 for a summary. In our study, the observed values for $\alpha$ are all approximately 1, as predicted by the \CB. Because of the low count statistics at high energies, we could not always fit $\beta$. We then fixed it to its theoretical value of 2.1 in order to make the fit converge and to obtain a value for the parameter $b$. In the cases where we could fit $\beta$, we found values close to 2.1 (Table~\ref{tab:CB_pars}). For the factor $b$ of the modified PL component in the \CB\ function we typically found values of the order $0.1\,$. An exception is GRB 031111, where $b$ is of the order 1.0, but with a large error ($0.4$). Our values for $\beta$ and $b$ are similar to the ones found by \citet{Dado2005} (fit of GRB 941017) and by \citet{Dado2004} (fit of $X$-ray flashes XRF 971019, XRF 980128, and XRF 990520 using BeppoSAX/WFC and CGRO/BATSE data). The authors of the \CB\ hypothesize that XRFs are simply GRBs viewed further off the jet axis. The different time development of the CPL- and the mPL-fluences, as reported in Table~\ref{tab:dt_grb021206}, possibly point to a different time dependence of the two underlying electron distributions within a cannonball. \subsection{Fitting \CB\ function versus CPL and Band function} Both the Band function and the \CB\ function are extensions of a CPL, the Band function with one additional parameter, the \CB\ function with two. For cases where a CPL fits the data well, also a Band function with $\beta = \infty$ or a \CB\ function with $b=0$ (and $\beta=2.1$ or any other value) fits. This is the case for GRB 031027. Whether additional parameters are necessary in a fit, can be tested with the $F$-test. For GRB 030329, the extra parameters are barely needed. For GRB 020715, 021008, 021206, 030406, 030519B, 031111 additional parameters are required at a confidence level of at least 90\%. Concerning the question of whether the high energy power law parameter $\beta$ in the \CB\ should be treated as free parameter, the answer is 'yes' from a theoretical point of view, but in practice, see Table~\ref{tab:chi2}, the improvements in $\chi^2$ are marginal or small for all bursts we studied. Our practice to freeze $\beta$ at its theoretically predicted value in cases of bad convergence seems to be acceptable. It is more difficult to compare the goodness of fit using the Band function compared to using the \CB\ function. The two functions are not independent, because they both are dominated by a CPL up to the peak energy and higher. In most cases of our study, the two functions fit the observed spectrum equally well with a slight preference for the Band function. At high energies however (typically above several times the peak energy) the two functions are different, the spectral hardening being a unique feature of the \CB\ function. There is only one case, namely GRB 021206, where this hardening is observed. For the rear data going up to high energies, a Band function fit gives $\chi^2 / \mbox{dof} = 133.3 / 74$ (see Table~\ref{tab:chi2}). This is not acceptable at $<0.01$\% probability of being accidentally so high. The \CB\ fit on the other hand gives $\chi^2 / \mbox{dof} = 82.7 / 73$, which is fully acceptable at a $20$\% level. We would like to stress again that, while the \CB\ gives acceptable fits for {\em all} cases, the Band function fails in one case. This seems enough to us to give some credit to the \CB. But it is, of course, no proof that the \CB\ is the only theory capable of describing the spectrum of GRB 021206. For example, a Band function plus a PL with $\gamma\approx 1.5$ would also fit. But there is no theory to predict such a shape. To our knowledge, \CB\ is to date the only existing GRB model % that explains the prompt GRB spectra from first principles. At this place we also would like to note the the mean $\alpha$-value found for the BATSE catalogue is 1 \citep[see][]{BATSEcatalog}). We cite from their summary: ``{\em We confirmed, using a much larger sample, that the most common value for the low-energy index is $\approx -1\,$}\footnote{ this corresponds to $+1$ in our notation} {\em \citep{Preece2000, Ghirlanda02}. The overall distribution of this parameter shows no clustering or distinct features at the values expected from various emission models, such as $-2/3$ for synchrotron \citep{Katz94, Tavani96}, $0$ for jitter radiation \citep{Medv00}, or $-3/2$ for cooling synchrotron \citep{GhisCel99}.}'' They do not mention the \CB\ which would explain $\alpha \approx 1$. Note that the $\beta$ values of the \CB\ are systematically lower than the $\beta$ values of the Band function, compare Tables \ref{tab:CB_pars} and \ref{tab:Band_pars}. From Band function fits to BATSE GRBs, it is known that $\beta$ is clustered around 2.3, with a long tail towards higher values, see \citet{BATSEcatalog}. For \CB\ we would expect $\beta$ to cluster at slightly lower values. For criticism of the \CB, see e.g.\ \citet{Hillas}, but see also the answer by \citet{Dar2006}. \subsection{The spectral hardening} The difference of a \CB\ spectrum and Band function is the hardening at high energies. This becomes visible---for the GRBs studied here---in the few \MeV\ region, but it depends on the peak energy and the factor $b$. For $\alpha=1.0$ and $b=0.10$ the hardening typically appears at several times the peak energy and the second term dominates at 10 times the peak energy. For the spectral fit of XRFs done by \citet{Dado2004}, the spectral hardening is expected in the few hundred \keV\ region, just where the number of photons detected runs low. Most of our GRBs also suffer from this lack of statistics at high energies, preventing the detection of a hardening. A spectral coverage of two decades and good detection efficiency at high energies is necessary to experimentally observe the full shape of the \CB\ function. In the case of GRB 021206 we were able to detect this hardening, thanks to RHESSI's broad energy range (30~\keV\ to 15~\MeV), and because this is one of the brightest GRBs ever observed. There is a GRB observed by SMM from 20 keV up to 100 MeV, namely GRB 840805. As reported by \citet{Share84}, the spectrum of this burst shows emission up 100 MeV. In order to fit the spectrum, ``a classical thermal synchrotron function plus a power law'' was used. The power law component was required to fit the data above about 6 MeV. This is a hint of a spectral hardening around 6 MeV, and we suppose that the spectrum of this GRB can be fit by a \CB\ function. The spectral hardening observed in GRB 941017 \citep{Nature2003} seems to be different. The photon index of GRB 941017 above a few \MeV\ is $\approx 1.0$. This case is discussed by \citet{Dado2005} as a possible additional feature in the \CB\ spectrum. \subsection{Outlook} \label{sec:outlook} In order to find more GRB spectra that show the hardening characteristic for the \CB\ function, strong GRBs have to be observed over a broad enough energy range. With the forthcoming GLAST mission, we expect that more such spectra will be observed. But also joint analyses with more than one instrument could reveal this hardening. We therefore suggest: $\bullet$ to search for \CB\ spectrum candidates among joint Swift/RHESSI GRBs and XRFs, and joint Swift/Konus GRBs. $\bullet$ to reanalyse some BATSE bursts. Looking at the BATSE spectra published by \citet{Ghirlanda2007}, we suppose that the \CB\ can possibly improve the fits of GRB 980329, GRB 990123, and GRB 990510. The same can be said for GRB 911031 as published by \citet{Ryde2006}. And GRB 000429, as published in Fig.\ 19 of \citet{BATSEcatalog}, looks like a promising candidate as well. $\bullet$ to search in KONUS data for suitable GRBs. $\bullet$ to add the \CB\ function to XSPEC in order to make it more accessible to the astronomical community. | 7 | 10 | 0710.2858 |
0710 | 0710.4153_arXiv.txt | It is shown that nearly-flat 3+1D spacetime emerging from a dual quantum field theory in 2+1D displays quantum fluctuations from classical Euclidean geometry on macroscopic scales. A covariant holographic mapping is assumed, where plane wave states with wavevector $\vec k$ on a 2D surface map onto classical null trajectories in the emergent third dimension at an angle $\vec\theta=l_P\vec k$ relative to the surface element normal, where $l_P$ denotes the Planck length. Null trajectories in the 3+1D world then display quantum uncertainty of angular orientation, with standard deviation $\Delta\theta=\sqrt{l_P/z}$ for longitudinal propagation distance $z$ in a given frame. The quantum complementarity of transverse position at macroscopically separated events along null trajectories corresponds to a geometry that is not completely classical, but displays observable holographic quantum noise. A statistical estimator of the fluctuations from Euclidean behavior is given for a simple thought experiment based on measured sides of triangles. The effect can be viewed as sampling noise due to the limited degrees of freedom of such a theory, consistent with covariant bounds on entropy. | It is not understood how a fundamentally quantum world creates the appearance, to internally constituted observers, of an approximately classical 3+1D classical spacetime, within which quantum fields operate. There are however indications that general relativity is in some sense ``holographic'' and that spacetime somehow emerges from a quantum theory with lower dimensionality. Examples exist where a conformal field theory with D spatial dimensions appears to be dual to a (supersymmetric) quantum field theory with gravity in D+1 spatial dimensions\cite{Maldacena:1997re,Witten:1998zw,Aharony:1999ti,Alishahiha:2005dj,Horowitz:2006ct}. Arguments based on black hole thermodynamics and evaporation\cite{Bekenstein:1972tm,Bardeen:gs,Bekenstein:1973ur,Bekenstein:1974ax,Hawking:1975sw,Hawking:1976ra,'tHooft:1985re,Susskind:1993if}, string theory\cite{Strominger:1996sh}, thermodynamics in nearly-flat space\cite{Jacobson:1995ab}, and classical relativity\cite{Padmanabhan:2006fn,Padmanabhan:2007en,Padmanabhan:2007xy, Padmanabhan:2007tm}, have led many authors to suspect that a region of 3+1D spacetime with gravity and the fields within it may have a dual description in terms of null surfaces swept out by a 2+1D quantum theory with a cutoff at the Planck length $l_P$ \cite{Padmanabhan:2006fn,Padmanabhan:2007en,Padmanabhan:2007xy, Padmanabhan:2007tm,'tHooft:1993gx,Susskind:1994vu,'tHooft:1999bw,Bigatti:1999dp,Bousso:2002ju}. This paper shows that such holographic theories generally display directly measurable quantum departures from classical general relativity. Recently, the author has argued\cite{Hogan:2007rz,Hogan:2007hc,Hogan:2007ci} that a nearly-flat holographic spacetime approaches classical behavior on macroscopic scales in a distinctive way that displays small but detectable quantum fluctuations from Euclidean geometry even at macroscopic separations. Specifically, the classical metric at widely separated points along any null trajectory displays a quantum complementarity of transverse position. Conversely, at small separations, angles become increasingly poorly defined until at the Planck scale the third dimension melts into the behavior of a 2+1D theory. This paper shows how this behavior arises simply from the geometry of the holographic mapping of the 2+1D states into 3+1D. This exercise helps to define the class of theories that have general relativity as an approximate classical limit, and that display the phenomena of holographic indeterminacy and noise. | 7 | 10 | 0710.4153 |
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0710 | 0710.3545_arXiv.txt | { For some time, the Draco dwarf spheroidal galaxy has garnered interest as a possible source for the indirect detection of dark matter. Its large mass-to-light ratio and relative proximity to the Earth provide favorable conditions for the production of detectable gamma rays from dark matter self-annihilation in its core. The Solar Tower Atmospheric Cherenkov Effect Experiment (STACEE) is an air-shower Cherenkov telescope located in Albuquerque, NM capable of detecting gamma rays at energies above 100 GeV. We present the results of the STACEE observations of Draco during the 2005-2006 observing season totaling 10 hours of livetime after cuts. } \begin{document} | Dark matter is thought to be an important component of the Universe and research into its nature is actively pursued using a variety of techniques. Dark matter may be weakly interacting massive particles (WIMPs) which would tend to accumulate at the bottom of gravitational potential wells, such as galaxies, where they could undergo self-annihilation processes. Depending on WIMP mass and branching ratios, a measurable flux of high energy gamma rays could result. The Draco dwarf spheroidal galaxy has long garnered interest as a potential source of concentrated dark matter~\cite{Tyler}. Draco has one of the highest known mass-to-light ratios ($M/L$), perhaps as high as $500 M_\odot/L_\odot$~\cite{Wu}. Current observations are consistent with a cuspy density profile~\cite{Lokas}, which would enhance the gamma-ray production rate. Furthermore, since Draco is a satellite of the Milky Way, its relative proximity to the Earth ($d \sim 75~kpc$)\cite{Bonanos} might allow for the detection of such gamma rays. | STACEE does not detect a significant signal from Draco, and sets upper limits on cross-sections for WIMP with rest-mass energy greater than about 150 GeV. {\bf Acknowledgments:} Many thanks go to the staff of the National Solar Tower Test Facility, who have made this work possible. This work was funded in part by the U.S. National Science Foundation, the Natural Sciences and Engineering Research Council of Canada, Fonds Quebecois de la Recherche sur la Nature et les Technologies, the Research Corporation, and the University of California at Los Angeles. | 7 | 10 | 0710.3545 |
0710 | 0710.4962_arXiv.txt | We have explored the galaxy disk/extended halo gas kinematic relationship using rotation curves (Keck/ESI) of ten intermediate redshift galaxies which were selected by {\ggkMgII} halo gas absorption observed in quasar spectra. Previous results of six edge--on galaxies, probed along their major axis, suggest that observed halo gas velocities are consistent with extended disk--like halo rotation at galactocentric distances of 25--72~kpc. Using our new sample, we demonstrate that the gas velocities are by and large not consistent with being directly coupled to the galaxy kinematics. Thus, mechanisms other than co--rotation dynamics (i.e., gas inflow, feedback, galaxy--galaxy interactions, etc.) must be invoked to account for the overall observed kinematics of the halo gas. \noindent In order to better understand the dynamic interaction of the galaxy/halo/cosmic web environment, we performed similar mock observations of galaxies and gaseous halos in $\Lambda$--CDM cosmological simulations. We discuss an example case of a $z=0.92$ galaxy with various orientations probing halo gas at a range of positions. The gas dynamics inferred using simulated quasar absorption lines are consistent with observational data. | 7 | 10 | 0710.4962 |
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0710 | 0710.1630_arXiv.txt | We derive new constraints on the Hubble function $H(\phi)$ and subsequently on the inflationary potential $V(\phi)$ from WMAP 3-year data combined with the Sloan Luminous Red Galaxy survey (SDSS-LRG), using a new methodology which appears to be more generic, conservative and model-independent than in most of the recent literature, since it depends neither on the slow-roll approximation for computing the primordial spectra, nor on any extrapolation scheme for the potential beyond the observable e-fold range, nor on additional assumptions about initial conditions for the inflaton velocity. This last feature represents the main improvement of this work, and is made possible by the reconstruction of $H(\phi)$ prior to $V(\phi)$. Our results only rely on the assumption that within the observable range, corresponding to $\sim$ 10 e-folds, inflation is not interrupted and the function $H(\phi)$ is smooth enough for being Taylor-expanded at order one, two or three. We conclude that the variety of potentials allowed by the data is still large. However, it is clear that the first two slow-roll parameters are really small while the validity of the slow-roll expansion beyond them is not established. | 7 | 10 | 0710.1630 |
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0710 | 0710.4934_arXiv.txt | {} {We present the results of the analysis of an archival observation of LMC X--2 performed with XMM/Newton. The spectra taken by high-precision instruments have never been analyzed before.} {We find an X--ray position for the source that is inconsistent with the one obtained by ROSAT, but in agreement with the Einstein position and that of the optical counterpart. The correlated spectral and timing behaviour of the source suggests that the source is probably in the normal branch of its X-ray color-color diagram. The spectrum of the source can be fitted with a blackbody with a temperature 1.5 keV plus a disk blackbody at 0.8 keV. Photoelectric absorption from neutral matter has an equivalent hydrogen column of $4 \times 10^{20}~\mathrm{cm}^{-2}$. An emission line, which we identify as the O VIII Lyman-$\alpha$ line, is detected, while no feature due to iron is detected in the spectrum.} {We argue that the emission of this source can be straightforwardly interpreted as a sum of the emission from a boundary layer between the NS and the disc and a blackbody component coming from the disc itself. Other canonical models that are used to fit Z-sources do not give a satisfactory fit to the data. The detection of the O VIII emission line (and the lack of detection of lines in the iron region) can be due to the fact that the source lies in the Large Magellanic Cloud. } {} | Low Mass X-ray Binaries (LMXBs) are binary systems harboring a compact object accreting mass from a late-type star. Systems hosting a weakly magnetized neutron star (NS) are usually divided into two classes: the Z sources, with luminosities close to the Eddington luminosity, $L_{\rm edd}$, and Atoll sources, usually with lower luminosities of $\sim 0.01-0.1\ L_{\rm edd}$. This classification relies upon the correlated X-ray spectral and timing properties, namely the pattern traced out by individual sources in the X-ray color-color diagram (CD, Hasinger \& van der Klis 1989). The seven known (Galactic) Z sources usually describe a complete Z-track in the CD on timescales of a few days, while Atoll sources cover their pattern on the CD on a longer timescale (usually several weeks). The correlated spectral and timing variability of LMC X--2 has suggested its classification as a Z-source (Smale \& Kuulkers 2000). An extensive study of these characteristics, using data taken with the PCA on board RXTE, has allowed to observe the complete Z-track in the CD which is completed in about 1 day (Smale, Homan \& Kuulkers 2003). Smale \& Kuulkers (2000) found evidence of a periodicity of $8.16~\mathrm{h}$ which they tentatively ascribe to an orbital period. A similar orbital period modulation has been found by Cornelisse et al. (2007) in the VLT optical lightcurve of the companion. From the shape of the Z-pattern described in the CD, it was inferred that LMC X-2 falls squarely into the category of the so-called ``Sco-like'' Z-sources due to its common characteristics with Sco X-1 (Smale et al. 2003). Its very high luminosity ($\sim 0.5-2 L_\mathrm{edd}$; Markert \& Clark 1975; Johnston, Bradt \& Doxsey 1979; Long, Helfand \& Grabelsky 1981; Bonnet-Bidaud et al. 1989; Christian \& Swank 1997; Smale \& Kuulkers 2000) makes it the brightest LMXB known together with Sco X--1. There are other analogies with Sco X--1: the optical counterpart is a faint, $M_V\sim 18.8$, blue star (similar to the optical counterpart of Sco X--1), and both sources show a similar correlation between the optical and the X-ray lightcurve during flares (McGowan et al. 2003). In some ways, we can consider LMC X--2 an extragalactic twin to Sco X--1. Measurement of Doppler effect on emission lines in the Bowen region allowed Cornelisse et al. (2007) to constrain the mass ratio of LMC X-2 to be $\le 0.4$. However, there is no recent X-ray spectral study of this source: a spectral analysis has been carried out by Bonnet-Bidaud et al. (1989) using EXOSAT data, where the spectrum of the source was fitted with a model consisting of a blackbody $kT\sim 1.3~\mathrm{keV}$ plus a thermal bremsstrahlung ($kT\sim 5~\mathrm{keV}$), or alternatively with a comptonized thermal model with $kT\sim 3~\mathrm{keV}$. They found no detectable feature in the Fe K-$\alpha$ range. In their survey of Low Mass X-Ray Binaries with Einstein, Christian \& Swank (1997) report that a fit in the 1-20 keV band with an unsaturated Comptonized model (i.e. a cutoff power-law) gave a reasonable fit with $\Gamma = 1.4$ and $kT = 6.4~\mathrm{keV}$. Schulz (1999), using ROSAT data in the 0.1--2.4 keV band, fitted the data using a blackbody with a temperature $kT = 1.5~\mathrm{keV}$ plus a thermal bresstrahlung model with a temperature of $kT=5~\mathrm{keV}$. Also, a gaussian emission line at 0.9 keV was evident in the spectrum. The accuracy of this study is obviously limited by the narrow band and the relatively poor spectral resolution of ROSAT. Smale \& Kuulkers (2000) studying RXTE/PCA data of the source found that the data can be well described by a cutoff power-law (e.g. a completely Comptonized component) with a cutoff temperature of $2.8~\mathrm{keV}$, or by a blackbody plus bremsstrahlung model with $kT_\mathrm{bb}\sim 1.5$ keV, $kT_\mathrm{brems} \sim 4.5$ keV. In this case the blackbody accounts for about 20\% of the total emission: as the authors point out, however, this model is not physically realistic given the enormous emitting volumes required for the bremsstrahlung emission. These low-resolution spectra differ noticeably from the X-ray spectra of other Z-sources, which are usually described in terms of a two-component model. The spectral models of Z-sources usually follow one of two main paradigms: the Eastern model (Mitsuda et al. 1989), in which the spectrum of the source is interpreted as a sum of an emission coming directly from the compact object (Comptonized or not) plus a blackbody emission from the disc, and the Western model (White et al. 1986), where the emission is due to a hot blackbody coming from the central source, plus a Comptonized component that is due to the emission from an accretion disc corona surrounding the disc (no direct thermal emission from the disc is seen as photons emitted there are reprocessed and Comptonized in the corona). Broad emission lines (FWHM up to $\sim 1$ keV) at energies in the range 6.4 -- 6.7 keV have been observed in the spectra of all galactic Z-sources, with the noticeable exception of GX~5-1 (Asai et al. 1994). These lines are identified with the K$\alpha$ radiative transitions of iron at different ionization states. Sometimes an iron absorption edge at energies $\sim 8$ keV has been detected (see e.g. Di Salvo et al. 2001). In this paper we analyze an archival XMM/Newton observation of LMC X--2, which gives us the opportunity to perform the first high resolution spectral study available to date on this source. | We have analyzed an archival XMM/Newton observation of LMC X--2, which allowed us to redetermine the position of the source with respect to the value reported in the Rosat All Sky Survey catalog. The source was probably in the normal or in the flaring branch, as its timing characteristics and high luminosity (exceeding the Eddington limit) exclude that the source is in the horizontal branch. Analyzing the spectra of the 3 active instruments we found that they could be well fitted by a two component model (blackbody plus disk blackbody) plus an emission line at 653 eV, which is identified as an O VIII Ly-$\alpha$ emission line. On the contrary, the iron region lacks any emission line. As discussed, this can be caused by the low metallicity of the stars in the LMC. | 7 | 10 | 0710.4934 |
0710 | 0710.4872_arXiv.txt | We have observed 13 methanol maser sources associated with massive star-forming regions; W3(OH), Mon~R2, S~255, W~33A, IRAS~18151$-$1208, G~24.78$+$0.08, G~29.95$-$0.02, IRAS~18556$+$0136, W~48, OH~43.8$-$0.1, ON~1, Cep~A and NGC~7538 at 6.7~GHz using the Japanese VLBI Network (JVN). Twelve of the thirteen sources were detected at our longest baseline of $\sim $50~M$\lambda $, and their images are presented. Seven of them are the first VLBI images at 6.7~GHz. This high detection rate and the small fringe spacing of $\sim$4~milli-arcsecond suggest that most of the methanol maser sources have compact structure. Given this compactness as well as the known properties of long-life and small internal-motion, this methanol maser line is suitable for astrometry with VLBI. | \label{section:introduction} The class II methanol masers are well-known as tracers of early stages of high-mass star formation (\cite{1998MNRAS.301..640W}; \cite{2001A&A...369..278M}; \cite{2006ApJ...638..241E}). Classes I and II are defined on the basis of associated sources (\cite{1987Natur.326...49B}; \cite{1991ASPC...16..119M}); different pumping mechanisms are proposed for them (\cite{1992MNRAS.259..203C}, \cite{1997MNRAS.288L..39S}). The class II masers are represented by 6.7 and 12.2~GHz lines (\cite{1991ApJ...380L..75M}; \cite{1987Natur.326...49B}). The spot size of class II masers is several AU (\cite{1992ApJ...401L..39M}; \cite{1999ApJ...519..244M}), while the class I masers are resolved out with Very Long Baseline Interferometric (VLBI) technique \citep{1998AAS...193.7101L}. VLBI observations for astrometry with hydroxyl, water and methanol masers have been made using the phase referencing technique. The accuracy of measuring with hydroxyl masers is $\sim$1~milli-arcsecond (mas) \citep{2003A&A...407..213V}, while that achieved with water masers is up to a few tens of micro-arcsecond ($\mu $as) \citep{2006ApJ...645..337H}. \citet{2006Sci...311...54X} also have achieved the positional accuracy of $\sim $10~$\mu $as with methanol maser at 12.2~GHz. This observation has been the only one astrometric observation with methanol masers. Internal proper motions are often measured typically at a few mas per year for water masers at 22~GHz (e.g., \cite{1981ApJ...247.1039G}, \yearcite{1981ApJ...244..884G}) and this motion sometimes prevents the separation of the annual parallax from the internal proper motion. The internal proper motion for class II methanol masers are small and have been measured only for W3(OH) at 12.2~GHz \citep{2002ApJ...564..813M}. The lifetime of 22~GHz water masers is sometimes too short for measuring annual parallax. For example, several spots disappear over timescales of 1 month for Cepheus~A and IRAS~21391$+$5802 (\cite{2001ApJ...560..853T}, \yearcite{2001Natur.411..277T}; \cite{2000ApJ...538..268P}). A monitoring of variability of 6.7~GHz methanol line for four years showed that each spectral feature survives regardless of some variability \citep{2004MNRAS.355..553G}, namely the lifetime of this maser is usually long enough for measuring annual parallax. There are 519 sites of 6.7~GHz methanol maser in a list compiled by \citet{2005A&A...432..737P}. Recently, new 48 sources have been detected using the 305~m Arecibo radio telescope \citep{2007ApJ...656..255P}. For the above reasons, i.e., compactness, small internal proper motion, long-life and large number of known sources, class II methanol masers may also be a useful probe for astrometry. The VLBI Exploration of Radio Astrometry (VERA; \cite{Kobayashi_etal.2003}) is a dedicated VLBI network for astrometry, which mainly observes galactic water masers for measuring their distance and motion. Methanol masers at 6.7~GHz would also contribute to revealing the Galactic structure as well as water masers. The spot size of methanol maser should be small enough for high precision astrometry. The study of spot size, however, have been made only a few cases so far. The number of sources imaged with baseline of $\geq $50~M$\lambda$ at 6.7~GHz were only four (\cite{1992ApJ...401L..39M}; \cite{2005A&A...442L..61B}; \cite{2006A&A...448L..57P}; \cite{2007A&A...461.1027G}). \citet{2002A&A...383..614M} discussed the size and structure of individual masing regions in detail based on their 12.2 and 6.7~GHz observations. They showed that the majority of the masing regions consist of a compact maser core surrounded by extended emission (halo), and derived the core size of 2 to 20~AU. We have started investigations of methanol masers using the Japanese VLBI Network (JVN) for astrometry. The network is a newly-established one with 50--2560~km baselines across the Japanese islands \citep{2006astro.ph.12528D} and consists of ten antennas, including four radio telescopes of the VERA. We have installed 6.7~GHz receivers on three telescopes of the JVN and have made a snap-shot imaging survey toward thirteen sources associated with massive star-forming regions. The aims of this observation are to investigate the detectability of 6.7~GHz masers with our longest baseline of $\sim$50~M$\lambda$ corresponding fringe spacing of $\sim$4~mas, and to verify the imaging capability of this network. In this paper, we describe the detail of this observation and data reduction in section \ref{section:observation}. In section \ref{section:result} we present the images of the detected sources and report the results on each individual source. Finally, we discuss the properties of this maser for astrometry in section \ref{section:discussion}, based on the results of this observation. \begin{table*} \caption{The sample of 6.7~GHz methanol masers}\label{tab:table1} \begin{center} \begin{tabular}{lllcrcc} \hline\hline Source & \multicolumn{2}{c}{Coordinates(J2000)} & Ref. & \multicolumn{1}{c}{$S_{\mathrm{p}}$} & $d$ & \multicolumn{1}{c}{VLBI obs.} \\ & \multicolumn{1}{c}{RA} & \multicolumn{1}{c}{Dec} & & & & \multicolumn{1}{c}{ } \\ & \multicolumn{1}{c}{(h m s)} & \multicolumn{1}{c}{($ ^{\circ} $ $ ^{\prime} $ $ ^{\prime\prime} $)} & & \multicolumn{1}{c}{(Jy)} & (kpc) & \\\hline W3(OH) & 02 27 03.820 & \hspace{1.5mm} 61 52 25.40 & 10 & 3294 & \hspace{0.8mm} 1.95 & 1 \\ Mon~R2 & 06 07 47.867 & $-$06 22 56.89 & 4 & 104 & \hspace{0.8mm} 0.83 & 6, 7 \\ S~255 & 06 12 54.024 & \hspace{1.5mm} 17 59 23.01 & 6 & 79 & 2.5 & 6, 7 \\ W~33A & 18 14 39.52 & $-$17 51 59.7 & 13 & 297 & 4.0 & \ldots \\ IRAS~18151$-$1208 & 18 17 58.07 & $-$12 07 27.2 & 13 & 119 & 3.0 & 8 \\ G~24.78$+$0.08 & 18 36 12.57 & $-$07 12 11.4 & \footnotemark[$*$] & 84 & 7.7 & \ldots \\ G~29.95$-$0.02 & 18 46 03.741 & $-$02 39 21.43 & 6 & 182 & 9.0 & \ldots \\ IRAS~18556$+$0136 & 18 58 13.1 & \hspace{1.5mm} 01 40 35 & 12 & 191 & 2.0 & \ldots \\ W~48 & 19 01 45.5 & \hspace{1.5mm} 01 13 28 & 3 & 733 & 3.4 & 6 \\ OH~43.8$-$0.1 & 19 11 53.987 & \hspace{1.5mm} 09 35 50.308 & 2 & 51 & 2.8 & \ldots \\ ON~1 & 20 10 09.1 & \hspace{1.5mm} 31 31 34 & 12 & 107 & 1.8 & \ldots \\ Cep~A & 22 56 17.903 & \hspace{1.5mm} 62 01 49.65 & 13 & 371 & \hspace{0.8mm} 0.73 & \ldots \\ NGC~7538 & 23 13 45.364 & \hspace{1.5mm} 61 28 10.55 & 6 & 256 & 2.8 & 5, 6, 7, 9, 11 \\\hline \multicolumn{7}{l}{\hbox to 0pt{\parbox{136mm}{\footnotesize Col~(1)~source name; Col~(2)--(3)~coordinates in J2000; Col~(4)~reference of coordinates; Col~(5)~peak flux densities from the single-dish observations by Yamaguchi 32~m; Col~(6)~source distance; Col~(7)~Published VLBI observations at 6.7~GHz. References --- (1)~\cite{1992ApJ...401L..39M}; (2)~\cite{1994ApJS...91..659K}; (3)~\cite{1995MNRAS.272...96C}; (4)~\cite{1998MNRAS.301..640W}; (5)~\cite{1998A&A...336L...5M}; (6)~\cite{2000A&A...362.1093M}; (7)~\cite{2001A&A...369..278M}; (8)~\cite{2002evn..conf..213V}; (9)~\cite{2004ApJ...603L.113P}; (10)~\cite{2005MNRAS.360.1162E}; (11)~\cite{2006A&A...448L..57P}; (12)~\cite{1994yCat.2125....0J}; (13)~Fringe rate mapping in our observations (accuracy of position is from 100 to 300~mas). \footnotemark[$*$] The coordinate of this source is not used for this data reduction. This coordinate is obtained from the observations made in the following year. It is coincident with the coordinate in the catalog listed by \citet{2005A&A...432..737P}.}\hss}} \end{tabular} \end{center} \end{table*} | \label{section:discussion} We have presented maps of twelve sources of methanol maser emission, and the seven of them are the first VLBI results at 6.7~GHz. The spatial distributions of maser spots show various morphology, such as linear (Cep~A, NGC~7538), ring like structure (W~48), largely separated clusters (W3(OH), W~33A, G~24.78$+$0.08, G~29.95$-$0.02, IRAS~18556$+$0136, OH~43.8$-$0.1, ON~1, Cep~A). We discuss the properties of methanol maser at 6.7~GHz as a possible probe for astrometry comparing with the methanol maser at 12.2~GHz and the water maser at 22~GHz. We have detected twelve out of thirteen methanol masers at 6.7~GHz with the longest baseline of 50~M$\lambda $ of our array. The integrated flux density of the correlated spectra account for $\sim$50~{\%} of the total flux, and some of which account for more than 90~{\%}. This high detection rate, flux recovery, and the small fringe spacing of 4~mas suggest that most of the methanol maser emission have compact structure. This result is consistent with a previous study for W3(OH) by \citet{1992ApJ...401L..39M}. We also showed that the correlated flux density at 50~M$\lambda$ is typically 20~{\%} of the total flux density. The size of maser spot inferred from the flux ratio varies from 2 to 30~AU (at a distance from 0.73 to 9~kpc of the sources). This is consistent with the core size of 2 to 20~AU obtained by \citet{2002A&A...383..614M}. The velocity range of typically 10~km~s$^{-1}$ \citep{1995MNRAS.272...96C} is narrow compared to that of water masers. Given this compactness and narrow velocity range as well as the known properties of long-life and small internal-motion, this methanol maser line is suitable for astrometry with VLBI. From the viewpoint of observation, atomospheric fluctuation is a significant problem for observations at higher frequency. Strong absorption by water vapor in the atomosphere makes observations relatively difficult at 22~GHz in summer. This would be a potential problem in measurement of annual parallax. This is not the case for 6.7~GHz observation. The astrometric VLBI observation uses continuum reference sources which are usually distant quasars. Such continuum sources show power-law, decreasing spectra, and the flux density is larger at 6.7~GHz than that of 12.2~GHz or 22~GHz. This property makes the astrometric observations easier in terms of detectability of reference sources. The 6.7~GHz line might have some disadvantages due to its lower frequency in comparison with the 12.2~GHz line: The interstellar scintillation broadens the size of maser and reference sources. The interstellar broadening is stronger at lower frequency, and might affect to measure the precise position of the sources. Also the ionospheric density fluctuation changes path-length, consequently affects phase measurement at lower frequency. These effects depend on frequency as $\nu^{-2}$, i.e., affect 3.3 times larger for 6.7~GHz than for 12.2~GHz. Although \citet{2002A&A...383..614M} discussed that the interstellar broadening is not significant at a scale larger that 1~mas, we have to take account of these effects for astrometry. We have made a phase-referencing VLBI observation with the JVN and achieved that the positional accuracy of $\sim$50~$\mu $as at 8.4~GHz and the image dynamic range of $\sim$50 on a target \citep{2006PASJ...58..777D}. The Bigradient Phase Referencing (BPR) was used for this observation, and a reference calibrator was separated by a \timeform{2D.1} from the target source. The following equation (\ref{equ:equ1}) can be used to derive an accuracy $\Delta \pi$ of annual parallax, \begin{eqnarray} \Delta \pi = \frac{\theta_{\mathrm{beam}}}{D \cdot \sqrt{N_{\mathrm{spot}} \cdot N_{\mathrm{obs}}}} \label{equ:equ1} \end{eqnarray} where $\theta_{\mathrm{beam}}$ is the minimum fringe spacing, $D$ is Dynamic range of the image, $N_{\mathrm{spot}}$ is the number of spots used in measuring annual parallax, and $N_{\mathrm{obs}}$ is the number of observations. If we make a series of observations for nine methanol maser spots for five epochs as in the cases of the observation by \citet{2006Sci...311...54X}, it is expected that an accuracy of $\sim$12~$\mu $as in annual parallax would be achieved. The flux density at 6.7~GHz is typically $\sim $10 times larger than that of 12.2~GHz \citep{1995MNRAS.274.1126C}, and the number of observable sources of 6.7~GHz is much larger than that of 12.2~GHz. This practical reason let us choose 6.7~GHz line than 12.2~GHz one. We have started to improve the JVN at 6.7~GHz in terms of sensitivity and the number of stations to detect weak masers and reference sources. Usuda 64~m is one of the newly participating telescopes. Assuming an aperture efficiency of 50~{\%} and $T_{\mathrm{sys}}$ of 50~K for all telescopes, it is expected that sensitivity for fringe detection (7~$\sigma $) would be less than 5~Jy. We showed that the correlated flux density at 50~M$\lambda $ is typically 20~{\%} of the total flux density. Hence, sources with total flux density of larger than 25~Jy would be potential targets for astrometric observation. The number of such sources found in the catalog by \citet{2005A&A...432..737P} is larger than 150. \bigskip The authors wish to thank the JVN team for observing assistance and support. The JVN project is led by the National Astronomical Observatory of Japan~(NAOJ) that is a branch of the National Institutes of Natural Sciences~(NINS), Hokkaido University, Gifu University, Yamaguchi University, and Kagoshima University, in cooperation with Geographical Survey Institute~(GSI), the Japan Aerospace Exploration Agency~(JAXA), and the National Institute of Information and Communications Technology~(NICT). | 7 | 10 | 0710.4872 |
0710 | 0710.1666_arXiv.txt | Although supernova explosions and stellar winds happens at scales bellow 100 pc, they affect the interstellar medium(ISM) and galaxy formation. We use cosmological N-body+Hydrodynamics simulations of galaxy formation, as well as simulations of the ISM to study the effect of stellar feedback on galactic scales. Stellar feedback maintains gas with temperatures above a million degrees. This gas fills bubbles, super-bubbles and chimneys. Our model of feedback, in which 10\%-30\% of the feedback energy is coming from runaway stars, reproduces this hot gas only if the resolution is better than 50 pc. This is 10 times better than the typical resolution in cosmological simulations of galaxy formation. Only with this resolution, the effect of stellar feedback in galaxy formation is resolved without any assumption about sub-resolution physics. Stellar feedback can regulate the formation of bulges and can shape the inner parts of the rotation curve. | 7 | 10 | 0710.1666 |
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0710 | 0710.3349_arXiv.txt | {\small After a period of inflationary expansion, the Universe reheated and reached full thermal equilibrium at the reheating temperature $\treh$. In this work we point out that, in the context of effective low-energy supersymmetric models, LHC measurements may allow one to determine $\treh$ as a function of the mass of the dark matter particle assumed to be either an axino or a gravitino. An upper bound on their mass and on $\treh$ may also be derived.} | \label{sect:intro} Dark matter (DM) remains an unknown component of the Universe. While it has so far escaped detection, its existence has been convincingly inferred from gravitational effects that it imparts on visible matter through rotational curves of spiral galaxies, gravitational lensing, etc,~\cite{susy-dm-reviews}. The effects of dark matter also can be seen on large structure formation and on anisotropy of the cosmic microwave background (CMB). The CMB in particular provides a powerful tool for determining the global abundance of cold DM (CDM). Recently, the Wilkinson Microwave Anisotropy Probe (WMAP) has performed a high-accuracy measurement of several cosmological parameters. In particular, the relic density of CDM has been determined to lie in the range~\cite{Spergel:2006hy} \beq \abundcdm =0.104 \pm 0.009. \label{Oh2WMAP} \eeq Since DM has to be electrically and (preferably) color-charge neutral, from the particle physics point of view, a natural candidate for DM is some weakly interacting massive particle (WIMP). Within standard cosmology, the WIMP is produced via a usual freeze-out mechanism from an expanding plasma. Among specific, well-motivated particle candidates for the WIMP, the by far most popular choice is a stable lightest neutralino $\chi$ of effective low-energy supersymmetry (SUSY) models. Most efforts have gone to exploring the lightest neutralino $\chi$ of the Minimal Supersymmetric Standard Model (MSSM) as the lightest supersymmetric particle (LSP) that, in the presence of R-parity, is stable and makes up all, or most of, the CDM in the Universe. In addition to an impressive experimental effort of direct and indirect searches for cosmic WIMPs, the Large Hadron Collider (LHC) will soon start exploring the TeV energy scale and is expected to find several superpartners and to determine their properties. In particular, some authors have explored the feasibility of determining the neutralino's relic abundance $\abundchi$ from LHC measurements~\cite{oh2atlhc-cmssm,oh2atlhc-mssm}. Their conclusion was that, under favorable circumstances, this should be possible with rather good accuracy, of order 10\% or better, although this may be challenging~\cite{bk05}. An analogous study has also been done in the context of the Linear Collider, where accuracy of a similar determination would be much better~\cite{oh2atilc}. If a WIMP signal is detected in one or more DM detection experiments, and also at the LHC a large missing-mass and missing-energy signature, characteristic of the stable neutralino, is measured and implies a similar mass, and if perhaps additionally eventually its relic abundance is determined from LHC data, even if with limited precision, and agrees with the ``WMAP range''~(\ref{Oh2WMAP}), then the DM problem will most likely be declared solved, and for a good reason. However, such an optimistic outcome is by no means guaranteed. One realistic possibility is that the neutralino, even if it is found as an apparently stable state in LHC detectors, may not be the true LSP and therefore DM in the Universe. Instead, it could decay in the early Universe into an even lighter, and (possibly much) more weakly interacting, state, the real LSP, outside the MSSM (or some other low-energy SUSY model) spectrum. In this case, current cosmic WIMP searches will prove futile, even after improving the upper limit on the spin-independent interaction cross section on a free proton $\sigsip$ from the current sensitivity of $\sim10^{-7}\pb$~\cite{silimit-07} down to $\sim10^{-10}\pb$, which is as far down as experiment can probably go given background from natural radioactivity. Moreover, the neutralino relic abundance, as determined at the LHC, may come out convincingly outside the range~(\ref{Oh2WMAP}). In fact, a value of $\abundchi$ above about 0.1 can be easily explained in terms of a lighter LSP into which the neutralino, after its freeze-out, decayed in the early Universe~\cite{ckr}. This is because the relic abundance of the true LSP is in this case related to $\abundchi$ by the ratio of the LSP mass $\mlsp$ to $\mchi$, which is less than one. If $\abundchi$ comes out below~(\ref{Oh2WMAP}), several solutions have been suggested which invoke non-standard cosmology, e.g. quintessence-driven kination~\cite{lowoh2-quint}, while preserving the neutralino as the DM in the Universe. However, if at the same time DM searches bring null results, this will provide a strong indication against the neutralino nature of DM. In contrast to the above attempts, we will consider a whole range of possible values of $\abundchi$ at the LHC, both below and above 0.1. We will work instead within the framework of standard cosmology but will not assume the neutralino to be the DM in the Universe. In this context we point an intriguing possibility of determining at the LHC the reheating temperature of the Universe. The framework we consider will therefore give us an opportunity to probe some crucial features of the early period of the Universe's history. The reheating temperature $\treh$ is normally thought of as the temperature at which, after a period of rapid inflationary expansion, the Universe reheated (or, more properly, defrosted), and the expanding plasma reached full thermal equilibrium. Determining $\treh$ cannot be done with the neutralino as DM since it freezes out at $\tf\simeq \mchi/24$, which is normally thought to be much below $\treh$.\footnote{The possibility of $\treh\lsim\tf$ has been explored in~\cite{gkr00} and more recently in~\cite{ggetal06}. In this case one could think of determining experimentally $\treh$ even with the neutralino as DM.} Here instead we assume a different candidate for the LSP (assuming R-parity) and cold DM, whose relic density depends, at least in part, on $\treh$. This is the case for either an axino or a gravitino. The spin-$1/2$ axino (the fermionic superpartner of an axion) and the spin-$3/2$ gravitino (the fermionic superpartner of a graviton) are both well-motivated. The former arises in SUSY extensions of models incorporating the Peccei-Quinn solution to the strong CP problem. The latter is an inherent ingredient of the particle spectrum of supergravity models. Both, like the axion, form a subclass of extremely weakly interacting massive particles (E-WIMPs)~\cite{Choi:2005vq}. The characteristic strength of their interactions with ordinary matter is strongly suppressed by a large mass scale, the Peccei-Quinn scale $\fa\sim 10^{11}\gev$ in the case of axinos and the (reduced) Planck scale $\mplanck\simeq 2.4\times 10^{18} \gev$ for gravitinos. The mass of the axino $\maxino$ is strongly model-dependent and can take values ranging from keV up to TeV~\cite{ckn}. The mass of the gravitino $\mgravitino$ in gravity-mediated SUSY breaking schemes is given by $\mgravitino\sim{\ms^2/\mplanck}$, where $\ms\sim10^{11}\gev$ is the scale of local SUSY-breaking in the hidden sector, and is expected in the GeV to TeV regime. In other schemes of SUSY breaking $\mgravitino$ can be (much) smaller. In this work we want to remain as model-independent as possible and will treat $\maxino$ and $\mgravitino$ as free parameters. Both axinos and gravitinos can be produced in decays of the neutralino (or another ordinary superpartner, e.g. the stau) after freeze-out, as mentioned above, or in thermal scatterings and decay processes of ordinary particles and sparticles in hot plasma at high enough reheating temperatures $\treh$ - hence their relic density dependence on $\treh$. The possibility of axinos in a KSVZ axion framework~\cite{ksvz} as cold DM was pointed out in~\cite{ckr,ckkr} and next studied in several papers~\cite{crs02,crrs04,bs04,rs06}, while axinos as warm DM was considered in~\cite{rtw}. The gravitino as a cosmological relic was extensively studied in the literature, starting from~\cite{gravitino-early,ekn84}, more recent papers on gravitino CDM include~\cite{mmy93,bbp98,bbb00,fengetal,rrc04,ccjrr}. For definiteness, in this work we will first assume the lightest neutralino to be the lightest ordinary superpartner and the next-to-lightest particle (NLSP), although below we will also consider the case of the stau. Our main result is that, assuming the axino or gravitino as the true LSP and CDM, and that the (apparently stable) neutralino is discovered at the LHC and its ``relic abundance'' $\abundchi$ is determined from LHC data, then one should be able to determine the reheating temperature in the Universe as a function of the LSP mass, or at least place an {\em upper bound} on it, as we show below. In the regime where thermal production dominates, in the axino case we find $\treh\propto \fa^2/\maxino$ while in the gravitino case $\treh\propto \mgravitino$. Alternatively, in the non thermal production dominated regime, we find an {\em upper bound} on the allowed mass range of a CDM particle. Furthermore, $\abundchi$ at the LHC can be expected to come out either below or above (or for that matter even accidentally agree with) the ``WMAP range''~(\ref{Oh2WMAP}). The same holds true also in the even more striking case of the lighter stau taking the role of the NLSP instead. In this case the very discovery of an (apparently stable) charged massive particle at the LHC will immediately imply that DM is made up of some state outside the usual spectrum of low-energy superpartners. In this case, dedicated studies of the differential photon spectrum may allow one to distinguish between the axino and the gravitino LSP~\cite{bchrr05}. We stress that, a detection of a (seemingly) stable neutralino at the LHC will not be sufficient to prove that the neutralino is the LSP and the CDM.\footnote{Note that, for the neutralino, or any other superpartner, to appear stable in an LHC detector, it is sufficient that its lifetime is longer than a microsecond or so. Of course, if it is unstable, then it will be immediately clear that the neutralino is not cosmologically relevant.} Even establishing an apparent agreement of $\abundchi$, as to be determined at the LHC, with the range~(\ref{Oh2WMAP}) will not necessarily imply the neutralino nature of DM, although admittedly it will be very persuasive. For this, a signal in DM searches must also be detected and the resulting WIMP mass must be consistent with LHC measurements of the neutralino mass. The paper is organized as follows. In section~\ref{production}, we review the thermal and non-thermal production of axinos and gravitinos and next explain our strategy for determining $\treh$ at the LHC. In sections~\ref{axino} and~\ref{gravitino}, we discuss the reheating temperature determination with axino LSP and gravitino LSP respectively, assuming the neutralino to be the NLSP. In section~\ref{staunlsp} we consider instead the lighter stau as the NLSP. We make final remarks and summarize our findings in section~\ref{summary}. | \label{summary} If the axino or the gravitino E-WIMP is the lightest superpartner and the dominant component of CDM in the Universe, then, under favorable conditions, it will be possible to determine the reheating temperature after inflation from LHC measurements. Basically, one will need to measure the neutralino mass and other parameters determining its ``relic abundance'' - a non-trivial and highly exciting possibility in itself. The quantity may be found to be different from the ``WMAP range''~(\ref{Oh2WMAP}), thus suggesting a non-standard CDM candidate. Additionally, one will need to know the mass of the gluino, either through a direct measurement, or via the gaugino unification mass relation, since it plays an important role in E-WIMP production at high $\treh$. The neutralino is likely to be sufficiently long-lived to appear stable in LHC detectors, but in the early Universe it would have decayed into the axino or the gravitino LSP. Since each of them is extremely weakly interacting, it would be hopeless to look for them in usual WIMP dark matter searches. In the axino LSP case, if thermal production dominates, a consistency of its relic abundance with the ``WMAP range''~(\ref{Oh2WMAP}) implies $\treh\propto \fa^2/\maxino$. If non-thermal production is dominant instead, then, for a known value of $\mchi$, the value of $\abundchi$ will imply an upper bound on $\maxino/\fa^2$, which in turn will imply a lower bound on $\treh$. For the gravitino LSP, in the thermal production dominated case the analogous dependence is $\treh\propto \mgravitino$. If non-thermal production dominates then the value of $\abundchi$ will imply an upper bound on $\mgravitino$. For both axino and gravitino one can express this as \beq \mlspmax=\min\{\mnlsp, \frac{\abundcdm}{\abundnlsp}\mnlsp\}, \label{eq:mlspmax} \eeq which is independent of $\treh$. For both relics, one can also derive an upper bound on $\treh$, although in the axino case has to further assume that it is cold DM. On the other hand, for both axino and gravitino, it will be difficult to distinguish between the TP and NTP regimes but some information may be helpful. Typically, when $\mchi$ is large and $\abundnlsp$ is also large then NTP is dominant, while if $\abundnlsp$ is small then TP typically dominates. Similar conclusions as for the neutralino (for both E-WIMPs) apply if the NLSP is the stau, instead. Additionally, the very discovery of a charged, massive and (seemingly) stable state at the LHC will immediately imply that dark matter in the Universe is not made up of the usual neutralino. With so many similarities between both E-WIMPs, a natural question arises whether it would be possible to distinguish them in a collider experiment. Unfortunately, with neutralino NLSP this is unlikely to be possible. On the other hand, with the stau NLSP being electrically charged, enough of them could possibly be accumulated. By comparing their 3-body decays, it may be possible to only tell whether Nature has selected the axino or the gravitino as the lightest superpartner, but to actually also determine their mass with some reasonable precision~\cite{Buchmuller:2004rq,bchrr05}. Needless to say, this would allow one to determine also the reheating temperature. In this work, we focused on demonstrating the principle of determining $\treh$ at the LHC, and did not concern ourselves with uncertainties of experimental measurements. These are likely to be substantial at the LHC, the case we have mostly focused on here, unless one considers specific models~\cite{bk05}. A more detailed study will be required to assess their impact. An analogous study can be performed in the context of the planned Linear Collider. We also implicitly assumed that only one relic dominates the CDM component of the Universe, which in principle does not have to be the case. For example, the axion and the axino could co-exist and contribute comparable fractions to the CDM relic density. (Note that this could be the case also in the ``standard'' case of the stable neutralino.) This would introduce an additional uncertainty and would lead to replacing the derived values of $\treh$ with an upper bound on the quantity. \medskip {\bf Acknowledgements} \\ This work is supported by an STFC grant. K.-Y.C. has been partially supported by a PPARC grant, by the Ministerio de Educacion y Ciencia of Spain under Proyecto Nacional FPA2006-05423 and by the Comunidad de Madrid under Proyecto HEPHACOS, Ayudas de I+D S-0505/ESP-0346. L.R is partially supported by the EC 6th Framework Programmes MRTN-CT-2004-503369 and MRTN-CT-2006-035505. R.RdA is supported by the program ``Juan de la Cierva'' of the Ministerio de Educaci\'{o}n y Ciencia of Spain. The author(s) would like to thank CERN for hospitality and support and the European Network of Theoretical Astroparticle Physics ENTApP ILIAS/N6 under contract number RII3-CT-2004-506222 for support. This project benefited from the CERN-ENTApP joint visitor's programme on dark matter, 5-9 March 2007 and the CERN Theory Institute ``LHC-Cosmology Interplay'', 25 June - 10 Aug 2007. \newcommand\jcap[3] % {{\it J.\ Cosmol.\ Astrop.\ Phys.\ }{\bf #1} (#2) #3} \newcommand\mnras[3] { % {{\it Mon.\ Not.\ R.\ Astron.\ Soc.\ }{\bf #1} (#2) #3}} \newcommand\apjs[3] { {{\it Astrophys.\ J.\ Supp.\ }{\bf #1} (#2) #3}} \newcommand\aipcp[3] { {{\it AIP Conf.\ Proc.\ }{\bf #1} (#2) #3}} | 7 | 10 | 0710.3349 |
0710 | 0710.5759_arXiv.txt | \noindent We present low resolution UV-blue spectroscopic observations of red giant stars in the globular cluster M53 ([Fe/H]$=-1.84$), obtained to study primordial abundance variations and deep mixing via the CN and CH absorption bands. The metallicity of M53 makes it an attractive target: a bimodal distribution of 3883 \AA\ CN bandstrength is common in moderate- and high-metallicity globular clusters ([Fe/H] $\geq -1.6$) but unusual in those of lower metallicity ([Fe/H] $\leq -2.0$). We find that M53 is an intermediate case, and has a broad but not strongly bimodal distribution of CN bandstrength, with CN and CH bandstrengths anticorrelated in the less-evolved stars. Like many other globular clusters, M53 also exhibits a general decline in CH bandstrength and [C/Fe] abundance with rising luminosity on the red giant branch. | Star-to-star variations of the strength of the 3883 \AA\ CN band are nearly universal in moderate-metallicity ([Fe/H] $\geq -1.6$) globular clusters of the Milky Way, but unusual in low-metallicity globular clusters ([Fe/H] $\leq -2.0$) (see Gratton, Sneden \& Carretta 2004 for a thorough review). Typically a CN bandstrength is measured as a magnitude difference between the flux within the 3883 \AA\ CN band and the flux in a nearby comparison band chosen to be free of CN absorption. The index $S(3839)$ introduced by Norris et al. (1981) is one such measure of CN band strength. Not only does $S(3839)$ vary from star to star within individual globular clusters, its distribution is often bimodal in moderate-metallicity ([Fe/H] $\geq -1.6$) clusters. It is often anticorrelated with the strength of CH absorption in the G band at 4300 \AA\ among giants of similar absolute magnitude. This anticorrelation indicates that CN-strong stars are depleted in carbon. In order to have strong CN bands with low [C/Fe] abundances, the CN-strong stars must also be enhanced in nitrogen. This abundance pattern has directed investigations into the origin of the bimodality toward stellar sources, where hydrogen shell burning can produce abundance enhancements and depletions such as those observed in globular clusters in C and N (and also O, Na, Mg, and Al). Meanwhile, few CN-strong giants are known in the halo field (see, e.g., Pilachowski et al. 1996; Shetrone et al. 1999), indicating that the enrichment process responsible for CN bimodality only operated in cluster (or proto-cluster) environments. This paper considers whether the apparent lack of CN bimodality in low-metallicity globular clusters is a true lack of abundance variation or simply an effect of low overall metal abundance. Previous investigations into the behavior of CN abundance at low metallicity have been inconclusive as to whether large variations in bandstrength can be present. Trefzger et al. (1983) and Shetrone et al. (1999) report several CN-strong red giants with large nitrogen abundances in M15 ([Fe/H]$=-2.26$), and Carbon et al. (1982) found variations of [C/Fe] and [N/Fe] but no strong 3883 \AA\ CN bands among the red giant stars of M92 ([Fe/H]$=-2.28$). Briley et al. (1994) showed that the carbon and nitrogen abundances measured within M92 would result in a bimodal distribution of $S(3839)$ if the overall metallicity were scaled up to [Fe/H]$=-1.0$ dex. Smith \& Norris (1982) and Briley et al. (1993) both studied M55 ([Fe/H]$=-1.81$, Harris 1996), and found that it is quite homogeneous in CN bandstrength. That lack of CN-strong stars could be the result of low overall metallicity inhibiting the formation of CN molecules, low [C/Fe] among the brighter red giants, or a sign that whatever enrichment process caused CN variations in moderate-metallicity globular clusters did not operate at lower metallicities. A study of additional globular clusters with metallicities similar to M55 could help decide between these alternatives. We chose M53 as the subject of this study because it is similar in metallicity to M55 (based on our isochrone fits discussed below), and will hopefully provide insight into whether low-metallicity globular clusters contain CN-enhanced stars. While our study does not comment directly on the process that causes abundance inhomogeneities in globular clusters, we note here that a lack of abundance variation in the lowest-metallicity clusters would provide a constraint on all hypotheses about that process. To understand the CN and CH index behavior, we compare spectroscopy of the 3883 \AA\ CN band for red giants in M53 to the Norris et al. (1981) study of NGC 6752, a moderate-metallicity globular cluster ([Fe/H]$=-1.56$) that shows clear CN bimodality. We also convert CH bandstrengths into [C/Fe] abundances using synthetic spectra and look at the relations between carbon abundance and CN bandstrengths. An intrinsic spread is found in CN bandstrength, but it is smaller than the range observed in NGC 6752. This result is interesting in the light of work by D'Antona et al. (2002), who use stellar evolution models to draw connections between helium abundance and horizontal branch morphology. They find that globular clusters with solar helium abundances have short, red horizontal branches, while extended blue horizontal branches result from the smaller core mass of helium-enhanced stars. In their model, the helium enrichment comes from AGB feedback, implying that globular clusters with red horizontal branches, like M53, should exhibit only minor CNO-product enrichment. \begin{figure} \plotone{f1.eps} \caption{ Color-magnitude diagram for the red giant branch of M53. The photometry is from Rey et al. (1998) for stars identified as red giants in Cuffey (1965). Filled circles are for stars observed with Kast, open circles are stars we did not observe, crosses are for stars dismissed in $\S$2 as nonmembers.} \end{figure} | 7 | 10 | 0710.5759 |
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0710 | 0710.5273_arXiv.txt | {It is expected that Be stars are surrounded by circumstellar envelopes, and that a significant fraction have companions. Achernar ($\alpha$\,Eri) is the nearest Be star, and is thus a favourable target to search for their signatures using high resolution imaging.} {We aim at detecting circumstellar material or companions around Achernar at distances of a few tens of AU.} {We obtained diffraction-limited thermal IR images of Achernar using the BURST mode of the VLT/VISIR instrument.} {The images obtained in the PAH1 band show a point-like secondary source located 0.280" north-west of Achernar. Its emission is 1.8\% of the flux of Achernar in this band, but is not detected in the PAH2, SiC and NeII bands. } {The flux from the detected secondary source is compatible with a late A spectral type main sequence companion to Achernar. The position angle of this source (almost aligned with the equatorial plane of Achernar) and its projected linear separation (12.3\,AU at the distance of Achernar) favor this interpretation.} | The southern star \object{Achernar} ($\alpha$\,Eridani, \object{HD 10144}) is the brightest and nearest of all Be stars (V $=0.46$ mag). Depending on the author (and the technique used) its spectral type ranges from B3-B4IIIe to B4Ve (e.g., Slettebak~\cite{slettebak82}, Balona et al.~\cite{balona87}). The estimated projected rotation velocity $v \sin i$ ranges from 220 to 270\,km.s$^{-1}$ and the effective temperature $T_{\rm eff}$ from $15\,000$ to $20\,000$\,K (see e.g., Vinicius et al.~\cite{vinicius06}, Rivinius \textit{priv. comm.}, Chauville et al.~\cite{chauville01}). It has been the subject of a renewed interest since its distorted photosphere was resolved by means of near-IR long-baseline interferometry (Domicano de Souza et al.~\cite{domiciano03}). Further interferometric observations revealed the polar wind ejected from the hot polar caps of the star (Kervella \& Domiciano~\cite{kervella06}), due to the von Zeipel effect (von Zeipel~\cite{vonzeipel24}). Our observations aim at studying two aspects of the close environment of Achernar: the circumstellar envelope at distances of up to a few tens of AU, and binarity. Mid-infrared (hereafter MIR) imaging is prefectly suited for these two objectives. Firstly, this wavelength range corresponds to where a circumstellar envelope becomes optically thick (asuming the emission is caused by free-free radiation, following for instance Panagia \& Felli~\cite{panagia75}). Secondly, the contrast between Achernar and a cool companion will be significantly reduced. We hereafter present the result of our imaging campaign with the VLT/VISIR instrument, to explore the environment of this star at angular distances of $\approx 0.1$ to 10". | We presented high resolution thermal infrared images of the close environment of Achernar. A point-like source is identified at an angular distance of 0.280" from Achernar, contributing for 1.8\% of the flux of Achernar in the PAH1 band. This source is not detected in our longer wavelength images, probably due to their insufficient sensitivity. Its location, close to Achernar and almost in the equatorial plane of the star, indicates that it is likely to be a physical companion orbiting Achernar. The measured flux corresponds to that of an A7V star (similar to Altair, for instance). A secure identification requires further observations, that will also determine if source B indeed shares the proper motion of Achernar ($\simeq 97$\,mas/yr). | 7 | 10 | 0710.5273 |
0710 | 0710.0984_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. Following the successful flight of the detector prototype (AMS-01) aboard the space shuttle, AMS-02 is expected to provide a significant improvement on the current knowledge of the elemental and isotopic composition of hadronic cosmic rays due to its long exposure time (minimum of 3 years) and large acceptance (0.5 m$^2$sr) which will enable it to collect a total statistics of more than $10^{10}$ nuclei. Detector capabilities for charge, velocity and mass identification, estimated from ion beam tests and detailed Monte Carlo simulations, are presented. Relevant issues in cosmic ray astrophysics addressed by AMS-02, including the test of cosmic ray propagation models, galactic confinement times and the influence of solar cycles on the local cosmic ray flux, are briefly discussed. } \FullConference{International Europhysics Conference on High Energy Physics\\ July 21st - 27th 2005\\ Lisboa, Portugal} \begin{document} | Data from AMS-02 will provide a major improvement with respect to existing results for the hadronic cosmic ray spectrum. A total statistics of more than $10^{10}$ events will be collected during its operation. Detector capabilities include charge separation up to $Z \sim 30$, velocity reconstruction with $\Delta \beta / \beta \sim 10^{-3}$ for $Z=1$ and $\Delta \beta / \beta \sim 10^{-4}$ for $Z \sim 10-20$, and isotopic separation of light elements. Results of AMS-02 will address key issues in cosmic ray astrophysics, namely, propagation of cosmic rays in the Galaxy, confinement times, and the effect of solar cycles on the flux in Earth's vicinity. | 7 | 10 | 0710.0984 |
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0710 | 0710.5429_arXiv.txt | We present $UBVI_C$ CCD photometry of the young open cluster Be 59 with the aim to study the star formation scenario in the cluster. The radial extent of the cluster is found to be $\sim$ 10 arcmin (2.9 pc). The interstellar extinction in the cluster region varies between $E(B-V) \simeq$ 1.4 to 1.8 mag. The ratio of total-to-selective extinction in the cluster region is estimated as $3.7\pm0.3$. The distance of the cluster is found to be $1.00\pm0.05$ kpc. Using near-infrared colours and slitless spectroscopy, we have identified young stellar objects (YSOs) in the open cluster Be 59 region. The ages of these YSOs range between $<1$ Myr to $\sim$ 2 Myr, whereas the mean age of the massive stars in the cluster region is found to be $\sim$ 2 Myr. There is evidence for second generation star formation outside the boundary of the cluster, which may be triggered by massive stars in the cluster. The slope of the initial mass function, $\Gamma$, in the mass range $2.5 < M/M_\odot \le 28$ is found to be $-1.01\pm0.11$ which is shallower than the Salpeter value (-1.35), whereas in the mass range $1.5 < M/M_\odot \le 2.5$ the slope is almost flat. The slope of the K-band luminosity function is estimated as $0.27\pm0.02$, which is smaller than the average value ($\sim$0.4) reported for young embedded clusters. Approximately $32\%$ of H$\alpha$ emission stars of Be 59 exhibit NIR excess indicating that inner disks of the T-Tauri star (TTS) population have not dissipated. The MSX and IRAS-HIRES images around the cluster region are also used to study the emission from unidentified infrared bands and to estimate the spatial distribution of optical depth of warm and cold interstellar dust. | High-mass star forming regions have been known for many years as OB associations and HII regions and they have been observed quite extensively on various aspects. However, the census of low mass stars in such regions has not been possible until recently. Recent advancement in detectors have permitted the detection of substantial population of low mass stars in OB associations. Recently relatively large number of low mass stars have been detected in a few OB associations (e.g. Upper Scorpios, the $\sigma$ and $\lambda$ Ori regions; Preibisch \& Zinnecker 1999, Dolan \& Mathieu 2002). Since this realization, surveys have demonstrated that the Initial Mass Functions (IMFs) are essentially the same in all star forming regions. The apparent difference may be due mainly to the inherent low percentage of high mass stars and the incomplete survey of low mass stars in high-mass star forming regions (e.g. Kroupa 2002, Preibisch \& Zinneker 1999, Hillenbrand 1997, Massey et al. 1995). The massive stars of OB associations have strong influence and significantly affect the entire star-forming region. Strong UV radiation from massive stars could evaporate nearby clouds and consequently terminate star formation. Herbig (1962) suggested that low and intermediate-mass stars form first and with the formation of most massive star in the region, the cloud gets disrupted and star formation ceases. Alternatively shock waves associated with the ionization front may squeeze molecular cloud and induce subsequent star formation. Elmegreen \& Lada (1977) proposed the expanding ionization fronts play a constructive role to incite a sequence of star-formation activities in the neighborhood. The question however is: could both of these effects occur in different parts of the same star-forming regions? Recent studies have shown cases where massive stars ionize adjacent clouds revealing embedded stars (e.g. the Eagle Nebula; Hester et al. 1996). Preibisch \& Zinnecker (1999) show that star formation in the upper Scorpius OB association was likely triggered by a nearby supernovae. Recently Lee et al. (2005) have found evidence for triggered star formation in the bright-rimmed clouds (BRCs) in the vicinity of O stars. Pandey et al. (2001, 2005) found that star formation in few young clusters may be a continuous process. For example, In the case of NGC 663, star formation seems to have taken place non-coevally in the sense that formation of low mass stars precedes the formation of most massive stars. Whereas, in the case of NGC 654 and NGC 3603, formation of low mass stars did not cease after the formation of most massive stars in the clusters. The Sharpless region S171 is a large HII region associated with the Cepheus OB4 stellar association (Yang \& Fukui 1992). The star cluster Be 59 is located at the center of the nebulous region with nine O7 - B3 stars. The physical properties of the ionized gas have been investigated in radio continuum and recombination lines. (e.g. Felli et al. 1977, Rossano et al. 1980, Harten et al. 1981). Yang \& Fukui (1992) have mapped the molecular gas distribution using the $J=1-0$ lines of $^{12}$CO and $^{13}$CO emission, indicating that Be 59 generates the ionization front on the surface of two dense molecular clouds. They suggested that the dense gas is in contact with the continuum source and ionization front is deriving shocks into the clumps. As the shock penetrates into the molecular clumps, it forms a compressed gas layer which is unstable against self-gravity (Elmegreen 1989, 1998). One can expect that a new generation of stars may form from the compressed gas layer. Thus the S171 region provides an opportunity to make a comprehensive exploration of the effects of massive stars on low mass star formation. Therefore in this paper, we present a multi-wavelength study of the star forming region S171. We have also carried out slitless spectroscopy to identify pre-main sequence H$\alpha$ emission stars in the region. In section 2 we present optical CCD photometric and slitless spectroscopic observations and brief description of data reduction. In sect 3, we discuss archival data set used in the present study. In the ensuing sections we discuss results and star formation scenario in the Be 59 region. | 7 | 10 | 0710.5429 |
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0710 | 0710.5103_arXiv.txt | {The dust component of the interstellar medium (ISM) has been extensively studied in the past decades. Late-type stars have been assumed as the main source of dust to the ISM, but recent observations show that supernova remnants may play a role on the ISM dust feedback.} {In this work, we study the importance of low and high mass stars, as well as their evolutionary phase, on the ISM dust feedback process. We also determine the changes on the obtained results considering different mass distribution functions and star formation history.} {We describe a semi-empirical calculation of the relative importance of each star at each evolutionary phase in the dust ejection to the ISM. We compare the obtained results for two stellar mass distribution functions, the classic Salpeter initial mass function and the present day mass function. We used the evolutionary track models for each stellar mass, and the empirical mass-loss rates and dust-to-gas ratio. } {We show that the relative contribution of each stellar mass depends on the used distribution. Ejecta from massive stars represent the most important objects for the ISM dust replenishment using the Salpeter IMF. On the other hand, for the present day mass function low and intermediate mass stars are dominant.} {We confirm that late-type giant and supergiant stars dominate the ISM dust feedback in our actual Galaxy, but this may not the case of galaxies experiencing high star formation rates, or at high redshifts. In those cases, SNe are dominant in the dust feedback process.} | RGB and AGB stars are known as the major continuous dust producers in the Universe, but also SN remnants have shown fast grain growth and dusty shells of M$_d$ $\sim 10^{-2} - 10^{1}$ M$_{\odot}$ \citep{noz03}. From the classical nucleation theory, the timescale for grain growth in the ISM can be estimated by $\tau_g = 4 s a (f n_i m_i v_i)^{-1}$, where $a$ is the dust mean size, $s$ is the material density, $f$ the sticking probability, $n_i$ the gas phase density, $m_i$ the atom mass and $v_i$ the velocity of the i-th atom to be added onto the grain surface. Considering typical ISM parameters, $a \sim 0.1 - 1 \ \mu$m, $s \sim 2$ g cm$^{-3}$, $n_i \sim 1$ cm$^{-3}$, $T \sim 10$ K, and assuming an efficiency $f = 0.1 - 1$, we find $\tau_g < 10^9$ yr. Dust particles are likely to be destroyed by shock waves \citep{draine79, mckee89}. Grain-grain or ion-grain collisions will lead to the shattering process, reducing or destroying dust particles. \citet{jones94, jones96} described the shattering process of ISM dust induced by SN blasts, obtaining destruction timescales of $\tau_{d} \sim 4 - 6 \times 10^8$ yr. Actually, as mentioned in these works, the accurate derivation of destruction timescales depends on the velocity of the shock waves, the frequency of SNe and the physical properties of the ISM, as density and temperature. Hence, the destruction timescales are shorter than the dust growth scale, the observed dust could not be explained by nucleation in loco, but had to be recently ($< 10^9$ yr) injected into the ISM \citep{tielens98}. It is commonly suggested that the ISM dust should be mostly originated from evolved low and intermediate mass stars but recent observations showed the presence of large quantities of dust ($M_d = 10^8$ M$_{\odot}$) in the early Universe (z $> 5$) \citep{hughes98,arc01,dunne03,bert03,maio04}. At that age low and intermediate mass stars were not evolved yet and, therefore, SNe are recognized as responsible for such material ejections \citep{tod01,mor03,sug06,dwek07}. In the post-shock phase of the SN remnant, the gas is generally cool and dense enough to allow dust formation and growth \citep{fal03,fal05}. Observationally, it is confirmed for SN1987A, which shows a $10^{-3}$ M$_{\odot}$ dust shell with a dust to gas ratio for heavier elements $\sim 0.3$. The same result was obtained for several other galactic and extragalactic SNe \citep{bar05,gomez07}. Therefore, SNe seem to play an important part on the ISM dust replenishment process \citep{dwek98}. However, if the short dust destruction timescale is taken into account, SNe could only be the main source of dust of the Galaxy if it presented a high star formation rates in its recent history. In this work we present a semi-empirical model, in which we study the role of different stellar mass functions in the output of dust ejected to the ISM. The model is described in Sect. 2, in Sect. 3 we show the main results and present a brief discussion, followed by the conclusions. | It is still unclear what is the main source of the dust feedback in our Galaxy. The dust must be formed in stars and ejected to the ISM and some authors have argued favoring cool late-type stellar winds, which present high mass-loss rates and dust is proved to be formed in these sites by observations. On the other hand, other plausible sources are the SNe ejecta. During the final evolutionary phases, high mass stars explode and supersonically eject a very rich gas to the ISM. Firstly, to determine the relative amount of dust ejected to the ISM for each stellar mass at each evolutionary phase we calculated the dust ejection during each evolutionary phase of the stars for different stellar mass distributions. As main result we showed that SNe are the main source of ISM dust feedback if a classic Salpeter IMF distribution is assumed. On the other hand, if we use the present day mass function, we show that the main sources of dust to the present Galaxy are the low and intermediate mass stars, representing more than 90\% of the total dust mass. Secondly, we studied the dust feedback process along the galactic time evolution, as done by \cite{mor03}, but including the effects of dust destruction by SN blasts. For simplicity we used a constant destruction rate, consistent with the current galactic physical parameters. During previous ages the SNe ejecta are dominant, in agreement with previous works. We showed that, considering the dust destruction, low and intermediate mass stars are dominant for a galactic age of $t > 10^9$yr, in a much higher proportion. The total dust mass of $\sim 10^7$ M$_{\odot}$ is obtained for a star formation rate of 5 M$_{\odot}$yr$^{-1}$. The dust destruction timescale depends on the SNe frequency, as well as the ISM density and temperature. It is probable that the destruction rate was higher earlier during the galactic evolution. In this case, the presented conclusions will stand, and the role of low and intermediate mass stars in later stages will be even higher. | 7 | 10 | 0710.5103 |
0710 | 0710.3934_arXiv.txt | A review of recent studies on a new mechanism of generation of large-scale magnetic field in a sheared turbulent plasma is presented. This mechanism is associated with the shear-current effect which is related to the ${\bf W} {\bf \times} {\bf J}$-term in the mean electromotive force. This effect causes the generation of the large-scale magnetic field even in a nonrotating and nonhelical homogeneous sheared turbulent convection whereby the $\alpha$ effect vanishes (where ${\bf W}$ is the mean vorticity due to the large-scale shear motions and ${\bf J}$ is the mean electric current). It is found that turbulent convection promotes the shear-current dynamo instability, i.e., the heat flux causes positive contribution to the shear-current effect. However, there is no dynamo action due to the shear-current effect for small hydrodynamic and magnetic Reynolds numbers even in a turbulent convection, if the spatial scaling for the turbulent correlation time is $\tau(k) \propto k^{-2}$, where $k$ is the small-scale wave number. We discuss here also the nonlinear mean-field dynamo due to the shear-current effect and take into account the transport of magnetic helicity as a dynamical nonlinearity. The magnetic helicity flux strongly affects the magnetic field dynamics in the nonlinear stage of the dynamo action. When the magnetic helicity flux is not small, the saturated level of the mean magnetic field is of the order of the equipartition field determined by the turbulent kinetic energy. The obtained results are important for elucidation of origin of the large-scale magnetic fields in astrophysical and cosmic sheared turbulent plasma. | \label{} Turbulence with a large-scale velocity shear is a universal feature in astrophysical plasmas. It has been recently recognized that in a sheared turbulent plasma with high hydrodynamic and magnetic Reynolds numbers a mean-field dynamo is possible even in a nonhelical and nonrotating homogeneous turbulence whereby a kinetic helicity and $\alpha$ effect vanish (see Rogachevskii and Kleeorin 2003; 2004; 2007; Brandenburg 2005; Brandenburg and Subramanian 2005c; Rogachevskii et al. 2006a; 2006b). The large-scale velocity shear produces anisotropy of turbulence with a nonzero background mean vorticity ${\bf W} = \bec{\nabla} {\bf \times} {\bf U}$, where ${\bf U}$ is the mean velocity. The dynamo instability in a sheared turbulent plasma is related to the ${\bf W} {\bf \times} {\bf J}$-term in the mean electromotive force, and it can be written in the form $\bec{\cal E}^\delta \propto - l_0^2 \, {\bf W} {\bf \times} (\bec{\nabla} {\bf \times} {\bf B}) \propto l_0^2 \, ({\bf W} {\bf \cdot} {\bf \Lambda}^B) {\bf B}$, where $l_0$ is the maximum scale of turbulent motions (the integral turbulent scale) and $ {\bf \Lambda}^B = \bec{\nabla} {\bf B}^2 / 2{\bf B}^2 $ determines the inhomogeneity of the mean original magnetic field ${\bf B}$. In a sheared turbulent plasma the deformations of the original magnetic field lines are caused by the upward and downward turbulent eddies, and the inhomogeneity of the original mean magnetic field in the shear-current dynamo breaks a symmetry between the influence of upward and downward turbulent eddies on the mean magnetic field. This creates the mean electric current ${\bf J}$ along the mean magnetic field and produces the mean-field dynamo due to the shear-current effect. The goal of this communication is to review recent studies on the new mechanism of generation of large-scale magnetic field due to the shear-current effect in a sheared turbulent plasma. The mean-field dynamo instability is saturated by the nonlinear effects. There are two types of the nonlinear effects caused by algebraic and dynamic nonlinearities. The effects of the mean magnetic field on the motion of fluid and on the cross-helicity result in quenching of the mean electromotive force which determines the algebraic nonlinearity. The dynamical nonlinearity in the mean-field dynamo is caused by the evolution of small-scale magnetic helicity, and it is of a great importance due to the conservation law for the magnetic helicity in turbulent plasma with very large magnetic Reynolds numbers (see, e.g., Kleeorin and Rogachevskii 1999; Brandenburg and Subramanian 2005a; Rogachevskii et al. 2006b, and references therein). The combined effect of the dynamic and algebraic nonlinearities saturates the growth of the mean magnetic field. The shear-current effect has been studied by Rogachevskii and Kleeorin (2003; 2004; 2007) for large hydrodynamic and magnetic Reynolds numbers using two different approaches: the spectral $\tau$ approximation (the third-order closure procedure) and the stochastic calculus (the path integral approach in a turbulence with a finite correlation time). A justification of the $\tau$ approximation for different situations has been performed in numerical simulations and analytical studies by Blackman and Field (2002); Field and Blackman (2002); Brandenburg at al. (2004); Brandenburg and Subramanian (2005a, 2005b); Sur et al. (2007). | 7 | 10 | 0710.3934 |
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0710 | 0710.4513_arXiv.txt | Extragalactic jets are visualized as dynamic erruptive events modelled by time-dependent magnetohydrodynamic (MHD) equations. The jet structure comes through the temporally self-similar solutions in two-dimensional axisymmetric spherical geometry. The two-dimensional magnetic field is solved in the finite plasma pressure regime, or finite $\beta$ regime, and it is described by an equation where plasma pressure plays the role of an eigenvalue. This allows a structure of magnetic lobes in space, among which the polar axis lobe is strongly peaked in intensity and collimated in angular spread comparing to the others. For this reason, the polar lobe overwhelmes the other lobes, and a jet structure arises in the polar direction naturally. Furthermore, within each magnetic lobe in space, there are small secondary regions with closed two-dimensional field lines embedded along this primary lobe. In these embedded magnetic toroids, plasma pressure and mass density are much higher accordingly. These are termed as secondary plasmoids. The magnetic field lines in these secondary plasmoids circle in alternating sequence such that adjacent plasmoids have opposite field lines. In particular, along the polar primary lobe, such periodic plasmoid structure happens to be compatible with radio observations where islands of high radio intensities are mapped. | Collimated jets with high terminal velocities appear to be universal phenomena in astrophysics. They are often associated with young stellar objects, compact galactic objects, and active galactic nuclei [Livio 1997]. These jets are always accompanied by accretion disks on the equatorial plane. The magnetosphere of the accretion disk contains a plasma that follows the accretion of the materials in the disk. The plasma density and pressure get higher as the accretion approaches to the center which drive the magnetic field. The dynamics of this system was first described in magnetohydrodynamic (MHD) model by Blandford and Payne [1982]. In this landmark paper, jets are considered as a spatial structure in stationary state. Ideal MHD equations in cylindrical coordinates $(r,\phi,z)$ are solved for time independent solutions. The central mass $M$ is replaced by a linear mass along the cylindrical axis. Self-similar solutions in space with a scale invariance $z/r$ are sought. Such self-similar solutions are compatible to a Keplerian disk plasma rotation velocity field superimposed by an Alfv\'{e}nic plasma velocity. The complete MHD instability spectrum with such Keplerian profile are analysed by Keppens, Casse, and Goedbloed [2002]. This model predicts that the magnetospheric disk plasma would be ejected towards the polar direction magnetocentrifugally should the magnetic field lines be at an angle less than $\pi/3$ or more that $2\pi/3$ on the $rz$ plane with respect to the outward radius of the disk. Collimating action of this plasma outflow would be provided by the hoop force of the azimuthal magnetic field and its associated parallel current in a force-free configuration. The interaction of the magnetic field with the plasma disk generates a MHD Poynting flux [Ferreira and Pelletier 1995] that can be converted into kinetic energy of the jet plasma [Zanni et. al. 2004]. Variants of jet formation model are proposed by Contopoulos and Lovelace [1994], Contopoulos [1995], and Cao and Spruit [1994]. This accretion-ejection model provides the basic framework of current investigations of accretion disk and jets as is reviewed by Balbus and Hawley [1998]. Dissipative MHD effects are examined by Casse and Ferreira [2000 a,b] and Casse [2004], and relativistic jets are analyzed by Vlahakis [2004]. These stationary state analytic studies are often complemented by numerical works in cylindrical geometry to simulate time evolutions of the jets [Ustyugova et. al. 1995, Ouyed and Pudritz 1997, Krasnopolski et. al. 1999], and the disk-jet system [Matsumoto et. al. 1996, Kato et. al. 2002]. Here, instead of a stationary model, we take the view that jets are a time dependent spatial structure. What we see is only a snapshot of their state at this particular moment. Due to their galactic dimensions, the time scale of these structures is believed to be extraordinarily large which gives the impression of a stationary structure. This implies that jets are results of an erruption originating from the galactic nucleus. They could be dissipated in time before another erruption takes place due to pressure built-up from accretion. Or one erruption could be superimposed on an earlier event. To model the jet system, we will do a self-similar analysis on the full time-dependent ideal MHD equations in spherical coordinates $(r,\theta,\phi)$ with a mass $M$ at the nucleus. In particular, we consider axisymmetric solutions. This type of self-similar solutions were pioneered by Low [1982a,b, 1984a,b] for astrophysical applications and solar corona mass ejections with pure radial plasma velocity flow. Variants of these solutions include cases where the plasma domain lies outside the mass $M$ such as interplanetary magnetic ropes in one-dimensional [Osherovich, Farrugia, and Burlaga 1993, 1995] and in two-dimensional [Tsui and Tavares 2005] cylindrical geometry, interplanetary magnetic clouds [Tsui 2006a], and also atmospheric ball lightnings [Tsui 2006b] in spherical geometry. In these descriptions where $M$ is outside the spherical domain of interest, the magnetic field is axisymmetric force-free and contains regions of closed field lines, while the plasma is spherically symmetric decoupled from the magnetic field. For the present case of extragalactic jets described by the mechanism of accretion-ejection, we will follow the self-similar solutions of Low with the polytropic index $\gamma=4/3$, but with particular emphasis on the finite plasma pressure. This current approach differs from the Keplerian disk plasma in that the radial flow is not tied to the Alfv\'{e}n velocity as a priori. In this dynamic model, we consider jets as a manifestation of mass ejection on a galactic scale. The plasma pressure, in this self-similar MHD model, proves to have an important role in collimating the magnetic fields and the jet plasmas. The time evolution function gives a dynamic description of the high radial flow especially in the jets. The self-similar solutions, that converge at the center and at infinity, give small regions of closed axisymmetric two-dimensional magnetic field lines where plasma density and pressure are much higher. These regions along the jets could correspond to the high intensity islands in radio frequency maps. \newpage | One of the main objections of self-organized plasmoid representation in free space in the absence of an adequate boundary is that it apparently violates the Virial theorem which states that [Schmidt 1966] \\ $${1\over 2}{d^2 I\over dt^2} +\int x_{k}{\partial G_{k}\over\partial t}dV\, =\,2(T+U)+W^{E}+W^{M} -\int x_{k}(P_{ik}+T_{ik})dS\,\,\,,\eqno(28)$$ \\ where $I$ is the moment of inertia of the plasmoid, $G$ is the momentum density of the electromagnetic field, $T$ and $U$ are the kinetic and thermal energies of the plasma, $W^{E}$ and $W^{M}$ are the electric and magnetic energies in the volume, $P_{ik}$ and $T_{ik}$ are the plasma and electromagnetic stress tensors. Taking the volume to cover the entire plasma and field, the surface term on the right side vanishes. In laboratory plasmas, this surface could be the machine vessel. Should the plasmoid be stationary, the volume term on the left side would be null, and the moment of inertia would be accelerating since the terms on the right side are all positive definite. While this statement has no conflict with the unbounded solutions, it apparently contradicts the stationary state of the bounded solutions. Nevertheless, this argument has overlooked the asymptotically bounded nature of the $H=-|H|<0$ plasmoid state. In this asymptotic case, we have $dI/dt=(dI/dy)(dy/dt)=0$ so that $I$ is stationary. The fact that $d^2I/dt^2>0$ implies that $I$ is at an asymptotic minimum, not an acceleration of $I$, which complies with the Virial theorem. The classical accretion-ejection model of Blandford and Payne [1982] is a time-independent stationary state model with spatial self-similar MHD solutions in cylindrical geometry $(r,\phi,z)$ with Alfv\'{e}nic plasma flow velocity plus a rotation, all with Keplerian scaling. In this model, the jets are formed by convecting the magnetospheric disk plasmas from the disk plane to the axis by magnetocentrifugal action through the magnetic field lines with low inclination angles to the disk plane. The angular momentum of the plasma on the disk plane is focused to the axis. The jets were put in place in the distant past according to the spatial self-similar solutions in $z/r$, and are maintained there by continuously transporting disk plasmas to the axis to sustain the axial outflow. Collimation to the axis is accomplished by magnetic hoop force. Here, we have taken a dynamic view where jets are the consequences of erruptive events, based on time-dependent MHD equations in spherical geometry $(r,\theta,\phi)$. The disk plasmas are accreted to the galactic nucleus where plasma pressure is built up to cause an erruption. The radially symmetric expanding velocity interacts the plasma with the magnetic field self-consistently through the MHD equations to generate spatial structures. Due to the existence of multiple quadratic invariants in the absence of dissipations, MHD systems have the tendency of developing self-organized and self-similar states through dissipative processes. The force-free configuration and the vortex nature of the magnetic field are akin to the quadratic invariants of the magnetic helicity and cross helicity of the MHD system. For these reasons, although temporally self-similar solutions are only a subset of general time-dependent MHD solutions, these self-similar configurations are prone to develope in natural phenomena. We, therefore, describe these spatial structures by self-similar temporal solutions in $\eta=r/y$. In this self-similar spherical model, consistent self-similar representations of plasma pressure, mass density, and magnetic fields are worked out for the $\gamma=4/3$ BC Low model, with special emphasis on the finite $\beta$ case. The spatial distribution of the magnetic field is described by an equation where the plasma pressure acts as the eigenvalue. The nature of this eigenvalue equation is as such that the magnetic field gets converged to the polar axis as a response to the high plasma pressure directly related to the eigenvalue. Since the separation constant $n(n+1)$ in this eigenvalue equation is a free parameter that as yet to be specified, there will be an adequate $n$ for almost any given plasma pressure. Although the radial plasma velocity is isotropic, the spatial structures are not. They could be highly collimated along the axial direction, and expand into the previously void space as time progresses. Although other types of solutions are permitted, we have specifically examined the solutions that bear resemblance with jet features with $n\gg 1$ and $a\eta_{0}\gg 1$. Since plasma and magnetic field are frozen into each other, plasma outflow is also collimated to the polar axis. The existence of regions of closed field lines along the primary polar magnetic lobe permits secondary plasmoids be embedded in it. These secondary plasmoids appear to be compatible to the observed islands of radio intensities. The time evolution function of the radial velocity consistent to the temporal self-similar solutions has different types of solutions according to the sign of $H$. Although the equatorial disk magnetospheric plasma is not addressed here, the accelerated accretion of plasma influx could be modelled with $H=0$ as the boundary condition at infinity. As for the polar jets, $H<0$ gives a bounded oscillating, or pulsating, solution. This bounded stage, we believe, nourishes the self-organized and self-similar states. With $H>0$, it gives an unbounded expanding solution with a high terminal velocity. It is apparent that plasma pressure due to accretion is the prime driving force that determines the value of $H$. The bounded oscillation would make a transition to the unbounded expansion as $H$ goes from negative to positive. According to our model, jet structures are, therefore, considered as a spatial configuration that has been expanding continuously into space. \newpage | 7 | 10 | 0710.4513 |
0710 | 0710.4039_arXiv.txt | {The X-ray source \source\/ is a very weak Low Mass X-ray Binary discovered in 1999 with \sax\/ and located in the Galactic Center. This region has been deeply investigated by the \integral\/ satellite with an unprecedented exposure time, giving us an unique opportunity to study the hard X-ray behavior also for weak objects. } {We analysed all available \integral\/ public and private Key-Program observations with the main aim of studying the long-term behavior of this Galactic bulge X-ray burster. } {The spectral results are based on the systematic analysis of all \integral\/ observations covering the source position performed between February 2003 and October 2006. \source\/ did not shows any flux variation along this period as well as compared to previous \sax\/ observation. Hence, to better constrain the physical parameters we combined both instrument data. } {Long \integral\/ monitoring reveals that this X-ray burster is a weak persistent source, displaying a X-ray spectrum extended to high energy and spending most of the time in a low luminosity hard state. The broad-band spectrum is well modeled with a simple Comptonized model with a seed photons temperature of $\sim$ 0.5 keV and an electron temperature of $\sim$24 keV. The low mass accretion rate ($\sim$ 2$\times$10$^{-10}$ M$_{sun}$/yr), the long bursts recurrence time, the small sizes of the region emitting the seed photons consisting with the inner disk radius and the high luminosity ratio in the 40--100 keV and 20--40 keV band, are all features common to the Ultra Compact source class. } { We report, for the first time, on the X-ray behavior of this source: observations with unprecedent wide energy band and sensitivity revealed this X-ray burster is a persistent source with an hard spectrum extending up to energies $>$ 100 keV. } | \source\/ was discovered in 1999 during the monitoring campaign of the Galactic Centre region performed by the \sax\//WFC (Cocchi et al. 1999 and in 't Zand et al. 1999). This instrument observed a type-I X-ray burst with exponential decays on August 1999, identifying the compact object as a weakly magnetized neutron star in a low mass X-ray binary system with a derived distance of 7 kpc (Cocchi et al. 2001). The source was followed-up by the \sax\//instruments to investigate the spectral properties: \source\/ shows a spectrum extended up to about 60 keV, with an intensity of $\sim$ 6 mCrab, both in the MECS (1--10 keV) and in the PDS (15--60 keV). The spectrum was fitted by an absorbed power law ($\Gamma\sim$2.2) and a column density N$_H$$\sim$2$\times$10$^{22}$cm$^{-2}$ or by a Comptonised emission model with electron temperature of $\sim$ 50 keV and a lower column density N$_H$$\sim$1$\times$10$^{22}$cm$^{-2}$ (Cocchi et al. 2001). The ROSAT source J171237.1-373834 is just 0\farcm6 from the \sax\/ position with a flux of 1.6 mCrab in the energy range 0.1--2.4 keV (in 't Zand et al. 1999). No optical counterpart has been identified yet within the 13\arcsec\ (1$\sigma$) ROSAT error circle radius. \source, monitored during the PCA/RXTE bulge scan program, is continuously active, apparently in two states: a slowly changing state, and a quicker one (in 't Zand et al. 2007). It was included in the IBIS/ISGRI soft Gamma-Ray survey catalog (Bird et al. 2007) as a transient X-ray buster with fluxes corresponding to $4.7\pm0.1$ mCrab and $4.0\pm0.2$ mCrab in the energy bands 20--40 keV and 40--100 keV, respectively. Recently, Chelovekov et al. (2006) reported two further burst detections with \integral\//IBIS in 2003-2004, both with the same peak flux. | \integral\/ observations reveal for the first time that this X-ray burster is a persistent source, displaying a very hard X-ray spectrum and spending most of the time in a low luminosity and hard state. A broad-band spectrum obtained combining the \integral\/ together with the \sax\/ data displays properties which fits well into the classification of low luminosity ($\sim$0.01 $ L_{\rm Edd}$) weakly magnetized neutron stars with "hard spectra'' (Barret et al. 2001): a broad-band spectrum extending up to 100 keV, dominated by a Comptonized component, with a optical depth of the Comptonizing corona of $\sim$3, seed photons temperature of $\sim$0.5 keV, an electron temperature of $\sim$24 keV. A few LMXBs appear to spend most of the time in this "hard state'' (Di Salvo \& Stella 2002), e.g. 4U 0614+091 (Piraino et al. 1999) or 4U1850-087 (Sidoli et al. 2006). The emission spectrum could be also fitted with a two component model consisting of a blackbody, or multi color blackbody component which represents emission from an optically thick accretion disk or from the neutron star surface, together with a Comptonized component which is interpreted as emission from a hot inner disk region or a boundary layer between the disk and the neutron star. We have estimate a black body radius of $\sim 5$ km for the black body component and the inner radius of the disk $< 11 km$ for the multi disk component. Although the data statistics does not allow us to distinguish between the two physical models, both values of the radius give an indication of the system compactness. We have also estimated the sizes of the emitting seed photons region, assuming their emission as blackbody with temperature T$_0$ and radius R$_{seed}$ following in 't Zand et al. (1999). We obtain $R_{\rm seed}\simeq 8.3\,{\rm km}\, (L_{\rm C}/10^{37}\,{\rm erg\, s}^{-1})^{1/2} (1+y)^{-1/2} (kT_0/1\,{\rm keV})^{-2}$, where $L_{\rm C}$ is the luminosity of the Comptonization component. Here, we have estimated the amplification factor of the seed luminosity by the Comptonization as $(1+y)$, where $y$ is the Comptonization parameter, $y=4kT_{\rm e}\max(\tau,\tau^2)/m_{\rm e} c^2$. The obtained values are $\sim$ 1 km and $\sim$ 10 km for the model with the blackbody component and with multicolor disk blackbody component, respectively. The parameters obtained using the latter model allows for a more reasonable interpretation. \begin{table*} \vspace{1cm} \caption{Results of the \source\/ fit, using different models: a simple comptonized model, a blackbody plus a Comptonized component and a multicolor disk blackbody plus a Comptonized component. Uncertainties are at the 90\% confidence level for a single parameter variation.} \label{fitsax} \vspace{0.3cm} \centering \begin{tabular}{ccccccccc} \hline\hline && & & & & &&\\ Model 1&{N$_{\rm H}$} &...&...& {T$_{\rm o}$}& {T$_{\rm e}$} &$\tau$ & $n_{\rm COMPTT}$ & $\chi^2$/d.o.f \\ &($10^{22}~cm^{-2}$)&...&... &($keV$) &(keV)&...&($10^{-3}$)&...\\ && & & & & & &\\ $comptt$&$1.3\pm0.1$&...&... &$0.47\pm0.02$ &$24^{+16}_{-10}$ &$2.7\pm1.3$ &$2.0^{+1.2}_{-0.8}$&142/134\\ && & & & & & &\\ \hline && & & & & & &\\ Model 2&{N$_{\rm H}$} & {T$_{\rm in}$}& ${\rm r_{in}({\cos}i)^{0.5}}$& {T$_{\rm o}$}& {T$_{\rm e}$} &$\tau$ & $n_{\rm COMPTT}$ & $\chi^2$/d.o.f \\ &($10^{22}~cm^{-2}$) &($keV$)&(km)&($keV$) &(keV)&...&($10^{-3}$)&...\\ && & & & & & &\\ $diskbb+comptt$&$1.4\pm0.3$&$0.7\pm0.3$ &$<11$ &$0.5\pm0.2$ &$24^{+22}_{-9}$ &$3\pm2$ &$1.6^{+1.6}_{-0.8}$&120/132\\ && & & & & & &\\ \hline && & & & & & &\\ Model 3 &{N$_{\rm H}$} & {T$_{\rm bb}$}& {R$_{\rm bb}$}& {T$_{\rm o}$}& {T$_{\rm e}$} &$\tau$ & $n_{\rm COMPTT}$ & $\chi^2$/d.o.f \\ &($10^{22}~cm^{-2}$) &($keV$)&(km)&($keV$) &(keV)&...&($10^{-3}$)&...\\ && & & & & & &\\ $bbody+comptt$&$1.2\pm0.2$&$0.6\pm0.1$&$5\pm1$&$1.2\pm0.2$ &$34^{+54}_{-14}$ &$2.0\pm1.5$ &$0.5^{+0.6}_{-0.2}$&119/132\\ && & & & & & &\\ \hline \hline \end{tabular}\\ \end{table*} These results show that the radius of the region providing the seed photons for the Comptonization emission is consistent with the inner radius of the accretion disk, showing that the seed photons are bounded in small region spatially coincident with the inner part of the disk. Moreover the seed photons temperature $T_0$ is consistent with the inner disk temperature $T_{in}$. By imposing $T_0$=$T_{in}$, the fit values correspond to a disk temperature of $0.57\pm0.08$ keV with a $\chi^2/dof$ of 121/133. This behavior suggests to identify the inner disk as the region providing the seed photons of the Comptonization emission. This result was observed in most of ultra compact X-ray source: 4U~1850-087, 4U~1820-303 and 4U~0513-40 (Sidoli et al. 2001), XTE J1751-305 (Gierli\'nski \& Poutanen 2005), XTE J1807-294 (Falanga et al. 2005). Assuming the persistent flux represents the average mass accretion rate and an accretion efficiency of $\eta = 0.2$ (corresponding, e.g., to $M_{\rm NS}=1.4 {\rm M}_{\sun}$ and $R_{\rm NS}=10$ km) and using multi color disk model luminosity, $L_{0.1-200 keV}=1.6\times10^{36} erg s^{-1}$, we get $\dot{M}\simeq 2.8 \times 10^{-10} {\rm M}_{\sun}$ yr$^{-1}$ for \source. This low mass accretion rate is consistent with the hard spectrum of \source\/, according to our present understanding on the X-ray bursters (for review see van der Klis 2006). Moreover, the burst recurrence time is dependent on the mass accretion rate on the neutron star: the faster new fuel is provided from the donor star, the shorter is the burst recurrence time. Long \integral\/ monitoring shows clearly that the frequency of bursts in this sources is low, as previously reported by in 't Zand et al. (2007), who estimated a recurrence time of 345--6507 hr. In 't Zand et al. (2007) have selected all 31 persistent X-ray bursters as reported in the \integral\/ survey (Bird et al. 2006) and demonstrated that the low mass accretion rates are accompanied by the hard X-ray spectra of the persistent emission. Almost all the Ultra Compact X-ray binaries have the highest values of the (40--100)keV/(20--40)keV hardness ratio. \source\/ was not included as supposed to have a transient nature. We now know it is a persistent and weak source, showing a high ratio L$_{(40-100)keV}$/L$_{(20-40)keV}$$\sim$0.90, in agreement with the value of 0.85$\pm$0.06 as derivated from recent data in Bird et al. 2007. Finally, the low mass accretion rate, the long bursts recurrence time, the small sizes of the region emitting the seed photons consisting with the inner disk radius and the high luminosity ratio in the 40--100 keV and 20--40 keV support the idea of In 't Zand et al. (2007) that this source could be a candidate Ultra Compact Binary. | 7 | 10 | 0710.4039 |
0710 | 0710.3105_arXiv.txt | We present HST/WFPC2 Linear Ramp Filter images of high surface brightness emission lines (either [OII], [OIII], or H$\alpha$+[NII]) in 80 3CR radio sources. We overlay the emission line images on high resolution VLA radio images (eight of which are new reductions of archival data) in order to examine the spatial relationship between the optical and radio emission. We confirm that the radio and optical emission line structures are consistent with weak alignment at low redshift ($z < 0.6$) except in the Compact Steep Spectrum (CSS) radio galaxies where both the radio source and the emission line nebulae are on galactic scales and strong alignment is seen at all redshifts. There are weak trends for the aligned emission line nebulae to be more luminous, and for the emission line nebula size to increase with redshift and/or radio power. The combination of these results suggests that there is a limited but real capacity for the radio source to influence the properties of the emission line nebulae at these low redshifts ($z < 0.6$). Our results are consistent with previous suggestions that both mechanical and radiant energy are responsible for generating alignment between the radio source and emission line gas. | \label{sec:Introduction} Radio galaxies are an important class of extragalactic objects: they represent one of the most energetic astrophysical phenomena; they may be used as probes of their environments; and they are unique probes of the early Universe \citep{McCarthy93}. As the 3CR sample of radio galaxies is radio flux-density selected \citep{Bennett62}, it provides an unbiased optical sample to study the host galaxies of these radio-loud active galactic nuclei (AGN). The sample has been well studied at multiple wavelengths and previous ground-based studies have investigated the spatial coincidence of optical emission nebulae and the radio source \citep{Baum88,Baum89a,McCarthy87,Chambers87}. \citet{Chambers87} and \citet{McCarthy87} first demonstrated the ``alignment effect'' at high redshift ($z \geqslant 0.6$) where the continuum optical and/or IR emission is aligned along the radio axis. An ``alignment effect'' for emission line gas was also shown by \citet{McCarthy87} and \citet{McCarthy95}. In this paper we focus on the alignment of the fine scale high surface brightness emission line gas imaged by the Hubble Space Telescope (HST) with the radio emission. The primary mechanisms for the emission line gas alignment are thought to be shocks induced by the radio jet and photoionization from the central AGN \citep[e.g.,][]{Baum89a,McCarthy93,Best00}. \citet{McCarthy89} demonstrated the redshift dependence of the alignment effect - at $z < 0.1$ there is no alignment, at $0.1 < z < 0.6$ some alignment is seen, and for $z > 0.6$ essentially all the powerful radio galaxies show strong alignment. \citet{Baum89b} found alignment at low $z$ if one compares the emission in the quadrants containing the radio source to quadrants without radio emission. Studies which attempt to break the redshift-radio power degeneracy suggest that the alignment depends on both redshift and radio power \citep[e.g.,][]{Inskip02b}. Thus, the alignment depends on both the presence of extended gas in the environment as well as the ability of the radio source to influence its environment. In this paper we examine the relationship of the high surface brightness emission line gas and radio emission over a wide range of redshift ($0.017 < z < 1.406$), focusing on the transition redshift range $z < 0.6$ where the alignment effect starts to turn on. Using the Wide Field and Planetary Camera 2 (WFPC2) on HST, we obtained images with resolutions of 0.$\arcsec$05 -- 0.$\arcsec$1, similar to or better than the synthesis imaging resolution obtained with the Very Large Array \footnote{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.} (VLA) at 5 GHz in the ``A'' configuration. The higher resolution allows us to probe denser gas in the center of these galaxies. In addition, the tunable Linear Ramp Filter (LRF) permits narrow band imaging of a selected spectral feature for a wide range of redshifts, enabling studies of a specified emission line for a sample such as the 3CR covering a large range of redshift. \citet{Biretta02} conducted an initial study using this data, and published a subset of the data. Here we expand on their analysis, presenting the complete set of detected images as well as statistical results. The paper is organized in the following fashion: we discuss the properties of the sample in $\S$\ref{sec:Sample}. $\S$\ref{sec:Observations} discusses the observations made and the data reduction methods used. In $\S$\ref{sec:Data Analysis} we outline the analysis of the data. $\S$\ref{sec:Properties} contains a discussion of our sources and the results of our study. In $\S$\ref{sec:Conclusions} we conclude by discussing the implications of our results. Appendix \ref{Appendix} contains notes for individual sources. | \label{sec:Conclusions} We present WFPC2 LRF images of selected emission lines in a sample of 3CR radio galaxies with bright emission lines. We detect clumpy, high surface brightness emission line gas with a wide range of morphologies. We examine the properties of the emission line gas and its relationship to the radio source, and in particular the nature of the alignment effect at $z < 0.6$ where the effect is beginning to become important. We find there is a weak alignment effect in our low $z$ sample ($z < 0.6$). The effect is clearly weaker than seen at high redshifts and is also somewhat weaker than seen at similar redshifts in ground-based data. This suggests that the alignment effect is dominated by the more extended lower surface brightness emission line gas which falls below our detection threshold in WFPC2. The alignment at low redshift is strong when the radio source and the nebula are both of order galactic scales (e.g., CSS sources). This implies there is mechanical energy input from the radio source as it propagates through the ambient gas in the host galaxy \citep[e.g.,][]{deVries99,Axon00}. There is a weak trend for the nebulae in the aligned sources (at all radio/optical size ratios) to have higher line luminosity. This suggests that the radio source preferentially provides energy to gas which is along the radio source axis. This is likely to be predominantly mechanical energy when the radio source and nebulae have similar sizes, and predominantly photoionization when the radio source is much larger than the nebula \citep[e.g.,][]{Best00,Inskip02a,Moy02,ODea02,Labiano05}. There is a weak trend for the size of the high surface brightness emission line nebulae to increase with increasing radio power and/or redshift over the range $z < 0.6$. This may be due to both the increasing presence of gas on large scales and the ability of the more powerful radio sources to provide additional mechanical energy and/or ionizing photons as a function of redshift. Thus, there is evidence for a modest excess of aligned sources at low redshift ($z < 0.6$). The alignment is strong when the radio source and emission line nebula are both on galactic scales. There are also weak trends for the aligned emission line nebulae to be more luminous, and for the emission line size to increase with redshift and radio power. The combination of these results suggests that there is a limited but real capacity for the radio source to influence the properties of the emission line nebulae at these low redshifts. | 7 | 10 | 0710.3105 |
0710 | 0710.3823.txt | We study the transition from inspiral to plunge in general relativity by computing gravitational waveforms of non-spinning, equal-mass black-hole binaries. We consider three sequences of simulations, starting with a quasi-circular inspiral completing 1.5, 2.3 and 9.6 orbits, respectively, prior to coalescence of the holes. For each sequence, the binding energy of the system is kept constant and the orbital angular momentum is progressively reduced, producing orbits of increasing eccentricity and eventually a head-on collision. We analyze in detail the radiation of energy and angular momentum in gravitational waves, the contribution of different multipolar components and the final spin of the remnant, comparing numerical predictions with the post-Newtonian approximation and with extrapolations of point-particle results. We find that the motion transitions from inspiral to plunge when the orbital angular momentum $L=L_{\rm crit}\simeq 0.8M^2$. For $L<L_{\rm crit}$ the radiated energy drops very rapidly. Orbits with $L\simeq L_{\rm crit}$ produce our largest dimensionless Kerr parameter for the remnant, $j=J/M^2\simeq 0.724\pm0.13$ (to be compared with the Kerr parameter $j\simeq 0.69$ resulting from quasi-circular inspirals). This value is in good agreement with the value of $0.72$ reported in \cite{Hinder2007}. These conclusions are quite insensitive to the initial separation of the holes, and they can be understood by extrapolating point particle results. Generalizing a model recently proposed by Buonanno, Kidder and Lehner \cite{Buonanno:2007sv} to eccentric binaries, we conjecture that (1) $j\simeq 0.724$ is close to the maximal Kerr parameter that can be obtained by any merger of non-spinning holes, and (2) no binary merger (even if the binary members are extremal Kerr black holes with spins aligned to the orbital angular momentum, and the inspiral is highly eccentric) can violate the cosmic censorship conjecture. | The research area of gravitational wave (GW) physics has reached a very exciting stage, both experimentally and theoretically. Earth-based laser-interferometric detectors, including LIGO \cite{Abramovici:1992ah}, GEO600 \cite{Luck:1997hv} and TAMA \cite{Ando:2001ej}, are collecting data at design sensitivity, searching for GWs in the frequency range $\sim 10-10^3$~Hz. VIRGO \cite{Caron:1997hu} should reach design sensitivity within one year, and the space-based interferometer LISA is expected to open an observational window at low frequencies ($\sim 10^{-4}-10^{-1}$~Hz) within the next decade \cite{Danzmann:1998}. The last two years have also seen a remarkable breakthrough in the simulation of the strongest expected GW sources, the inspiral and coalescence of black-hole binaries \cite{Pretorius2005a,Campanelli2006, Baker2006}. Several groups have now generated independent numerical codes for such simulations \cite{Pretorius2006, Baker2006a, Campanelli2006a, Sperhake2006, Scheel2006, Bruegmann2006a, Herrmann2007, Pollney2007, Etienne2007} and studied various aspects of binary black hole mergers. % including the resulting gravitational recoil %\cite{Herrmann2007,Gonzalez2007,Koppitz2007,Gonzalez2007a,Campanelli2007v2, % Baker2007,Tichy2007,Herrmann2007c,Schnittman:2007ij,Lousto:2007db,Sopuerta:2006wj,Sopuerta:2006et} %spin precession and spin flips \cite{Campanelli2007b,Campanelli2007v2}, In the context of analyzing the resulting gravitational waveforms, these include in particular the comparisons of numerical results with post-Newtonian (PN) predictions \cite{Baker2006c, Schnittman2007, Hannam2007, Buonanno2006, Berti:2007fi, Boyle2007}, multipolar analyses of the emitted radiation \cite{Buonanno2006, Berti:2007fi, Schnittman:2007ij}, the use of numerical waveforms in data analysis \cite{Ajith:2007qp, Pan2007, Vaishnav2007, Ajith:2007kx} and gravitational wave emission from systems of three black holes \cite{Campanelli:2007ea}. Despite this progress, a comprehensive analysis of binary black hole inspirals remains a daunting task, mainly because of the large dimensionality of the parameter space. In geometrical units, the total mass of the binary is just an overall scale factor. The source parameters to be explored by numerical simulations (sometimes called ``intrinsic'' parameters in the GW data analysis literature) include the mass ratio $q=M_2/M_1$, the eccentricity $e$ of the orbit and six parameters for the magnitude of the individual black hole spins and their direction with respect to the binary's orbital angular momentum. In this paper we present results from numerical simulations of non-spinning, equal-mass black-hole binaries, and we focus on the effect of the orbital eccentricity on the merger waveforms. We consider three sequences, starting with quasi-circular inspirals that complete $\sim 1.5$, $\sim 2.3$ and $\sim 9.6$ orbits, respectively, prior to coalescence of the holes. By fixing the binding energy of the system and progressively reducing the orbital angular momentum, we produce a sequence of orbits of increasing eccentricity and eventually a head-on collision. For each of these simulations we analyze in detail the radiation of energy and angular momentum in GWs, the contribution of different multipolar components and the final spin of the remnant, comparing numerical predictions with the PN approximation and with extrapolations of point-particle results. Non-eccentric inspirals are usually considered the most interesting cases for GW detection. For an isolated binary evolving under the effect of gravitational radiation reaction, the eccentricity decreases by roughly a factor of 3 when the orbital semimajor axis is halved \cite{Peters:1963ux}. For most conceivable formation mechanisms of solar-mass black hole binaries, the orbit will usually be circular by the time the GW signal enters the best-sensitivity bandwidth of Earth-based interferometers. % However, we wish to stress that our simulations could be of interest for GW detection. For example, according to some astrophysical scenarios, eccentric binaries may be potential GW sources for Earth-based detectors. In globular clusters, the inner binaries of hierarchical triplets undergoing Kozai oscillations can merge under gravitational radiation reaction, and $\sim 30\%$ of these systems can have eccentricity $\sim 0.1$ when GWs enter the detectors' most sensitive bandwidth at $\sim 10$~Hz \cite{Wen:2002km}. %In globular clusters, binaries with essentially a thermal distribution of %eccentricities may exist \cite{Benacquista2002}. Massive black hole binaries to be observed by LISA could also have significant eccentricity in the last year of inspiral. %Analytical calculations and $N$-body simulations show that, in purely %collisionless spherical backgrounds, the expected equilibrium distribution of %eccentricities is skewed towards high $e\simeq 0.6-0.7$, and that dynamical %friction does not play a major role in modifying such a distribution %(\cite{1999ApJ...525..720C}, in particular Fig.~5). Recent simulations using smoothed particle hydrodynamics follow the dynamics of binary black holes in massive, rotationally supported circumnuclear discs \cite{2006MNRAS.367..103D, 2007MNRAS.379..956D, Mayer:2007vk}. In these simulations, a primary black hole is placed at the center of the disc and a secondary black hole is set initially on an eccentric orbit in the disc plane. By using the particle splitting technique, the most recent simulations follow the binary's orbital decay down to distances $\sim 0.1$~pc. Dynamical friction is found to circularize the orbit if the binary {\it corotates} with the disc \cite{2007MNRAS.379..956D}. However, if the orbit is {\it counterrotating} with the disc the initial eccentricity does not seem to decrease, and black holes may still enter the GW emission phase with high eccentricity \cite{2006MNRAS.367..103D}. Complementary studies show that eccentricity evolution may still occur in later stages of the binary's life, because of close encounters with single stars and/or gas-dynamical processes. Three-body encounters with background stars have been studied mainly in spherical backgrounds. These studies find that stellar dynamical hardening can lead to an increase of the eccentricity, acting against the circularization driven by the large-scale action of the gaseous and/or stellar disc, possibly leaving the binary with non-zero eccentricity when gravitational radiation reaction becomes dominant \cite{1996NewA....1...35Q,Aarseth:2002ie,Berczik:2005ff,Berczik:2006tz,Matsubayashi:2005eg}. It has also been suggested that the gravitational interaction of a binary with a circumbinary gas disc could increase the binary's eccentricity. The transition between disc-driven and GW-driven inspiral can occur at small enough radii that a small but significant eccentricity survives, typical values being $e\sim 0.02$ (with a lower limit $e\simeq 0.01$) one year prior to merger (cf. Fig.~5 of \cite{Armitage:2005xq}). If the binary has an ``extreme'' mass ratio $q\lesssim 0.02$ the residual eccentricity predicted by this scenario can be considerably larger, $e\gtrsim 0.1$. Numerical simulations should be able to test these predictions in the near future. As shown by Sopuerta, Yunes and Laguna, eccentricity could significantly increase the recoil velocity resulting from the merger of non-spinning black-hole binaries \cite{Sopuerta:2006et}. Independently of the presence of eccentricity in astrophysical binary mergers, the problem we consider here has considerable theoretical interest. Our simulations explore the transition between gravitational radiation from a quasi-circular inspiral (the expected final outcome in most astrophysical scenarios) and the radiation emitted by a head-on collision, where the binary has maximal symmetry. Our work should provide some guidance for analytical studies of the ``transition from inspiral to plunge''. The first analytical study of this problem in the context of PN theory was carried out by Kidder, Will and Wiseman \cite{Kidder:1993}. The transition between the adiabatic phase and the plunge was studied in \cite{Buonanno2000} using nonperturbative resummed estimates of the damping and conservative parts of the two-body dynamics, i.e. the so-called ``Effective One Body'' (EOB) model. Ori and Thorne \cite{Ori:2000zn} provided a semi-analytical treatment of the transition in the extreme mass ratio limit. Waveforms comprising inspiral, merger and ringdown for comparable-mass bodies have also been produced using the EOB model (see eg.~\cite{Buonanno:2005xu} for extensions of the original model to spinning binaries and for references to previous work). Preliminary comparisons of EOB and numerical relativity waveforms showed that improved models of ringdown excitation \cite{Buonanno2006,Pan2007,Berti:2006wq} or additional phenomenological terms in the EOB effective potential \cite{Buonanno:2007pf} are needed to achieve acceptable phase differences between the numerical and analytical waveforms. Our study is complementary to Ref.~\cite{Pretorius:2007jn}, that considered sequences of eccentric, equal-mass, non-spinning binary black hole evolutions around the ``threshold of immediate merger'': a region of parameter space separating binaries that quickly merge to form a final Kerr black hole from those that do not merge in a short time. Similar scenarios have also been studied in Ref.~\cite{Washik2008}, with particular regard to the maximal spin of the final hole generated in this way. The universality of the gravitational wave signal during the merger was analysed in Ref.~\cite{Hinder2007}, where it was pointed out that binaries largely circularize after about 9 orbits when starting with eccentricities below about $0.4$. The first comparison between numerical evolutions of eccentric binaries with post-Newtonian predictions was presented in \cite{Hinder2008}. Our focus in this work is on the high-eccentricity region of the parameter space, which always leads to merger. In particular, the near-head-on limit of our study is of interest as a first step to compute the energy loss and production cross-section of mini-black holes in TeV-scale gravity scenarios (possibly at the upcoming LHC \cite{Giddings:2007nr}), and trans-Planckian scattering in general \cite{Dray:1984ha,Giddings:2007bw}. Present semi-analytical techniques (including a trapped surface search in the union of Aichelburg-Sexl shock waves, close-limit approximation calculations and perturbation theory) only give rough estimates of the emitted energy and production cross-section \cite{Cardoso:2005jq} and do not provide much insight into the details of the process (but see \cite{Sperhake:2008ga} for a first numerical investigation). Our main finding is that, for all sequences we studied, the motion radically changes character when the black holes' % initial momentum $P\sim P_{\rm crit}\simeq 0.08-0.09M$ and the orbital angular momentum $L\sim L_{\rm crit}\simeq 0.8 M^2$, turning from an eccentric inspiral into a plunge. In particular, for $L\lesssim L_{\rm crit}$ we observe that: \begin{itemize} \item The number of orbits $N_{\rm waves}$ (as estimated using the gravitational wave cycles) or $N_{\rm punc}$ (as computed from the punctures' trajectories) becomes less than one, so the motion effectively turns into a plunge (see Table \ref{tab: models} and Fig.~\ref{fig: traj} below); \item The energy emission starts decreasing exponentially (Fig.~\ref{fig: El}); \item PN-based eccentricity estimates yield meaningless results (Table \ref{tab: models}); \item The polarization becomes linear rather than circular (Fig.~\ref{fig: pol}); \item The final angular momentum starts {\it decreasing}, rather than increasing, as $P$ and $L$ decrease (Fig.~\ref{fig: jfin}). \end{itemize} Binary mergers with $L\simeq L_{\rm crit}$ are those producing the largest Kerr parameter for the final black hole observed in our simulations, $j_{\rm fin}\simeq 0.724$. One is led to suspect that for maximally spinning holes having spins aligned with the orbital angular momentum, a large orbital eccentricity may lead to violations of the cosmic censorship conjecture. Using arguments based on the extrapolation of point-particle results (see also \cite{Buonanno:2007sv}), we conjecture that (1) the maximal Kerr parameter that can be obtained by any merger of non-spinning holes is not much larger than $j\simeq 0.724$, and (2) cosmic censorship will {\it not} be violated as a result of any merger, even in the presence of orbital eccentricity. Further numerical simulations are needed to confirm or disprove these conjectures. The paper is organized as follows. We begin in Sec.~\ref{eccentricity} discussing to what extent the Newtonian concept of eccentricity can be generalized to characterize orbiting binaries in general relativity. For this purpose, we introduce and compare various PN estimates of the orbital eccentricity, and we show that these eccentricity estimates break down when the motion turns from inspiral to plunge. Sec.~\ref{numerics} is a brief introduction to the numerical code used for the simulations. After a discussion of the choice of initial data and of the code's accuracy, we show how reducing the orbital angular momentum affects the gravitational waveforms, the puncture trajectories and the polarization of the waves. In Sec.~\ref{energyj} we study the multipolar energy distribution of the radiation and the angular momentum of the final Kerr black hole. In Sec.~\ref{BKL} we show that the salient features of our simulations can be understood using extrapolations of point-particle results. Sec.~\ref{QNMs} is devoted to fits of the ringdown waveform and to estimates of the energy radiated in ringdown waves. We conclude by considering possible future extensions of our investigation. %\clearpage %============================================================================= | We have presented a study of the gravitational waveforms produced by sequences of equal-mass, non-spinning black hole binaries. For each sequence, the binding energy of the system is kept constant and the orbital angular momentum is progressively reduced to zero, producing orbits of increasing eccentricity and eventually a head-on collision. We find that the motion transitions from inspiral to plunge when the orbital angular momentum $L=L_{\rm crit}\simeq 0.8M$. For $L<L_{\rm crit}$ the binary always completes less than $\sim 1$ orbit, and PN estimates of the orbital eccentricity are no longer meaningful. As the initial momentum of the holes $P/M\to 0$ the polarization quickly becomes linear, rather than circular, and the radiated energy drops (roughly) exponentially. For equal-mass, non-spinning binaries, orbits with $L\simeq L_{\rm crit}$ produce the largest dimensionless Kerr parameter for the remnant, $j_{\rm fin}=J/M^2\simeq 0.724\pm0.13$ (to be compared with the Kerr parameter $j_{\rm fin}\simeq 0.69$ resulting from quasi-circular inspirals). These results are quite insensitive to the initial separation of the holes, and they can be understood using extrapolations from black hole perturbation theory. % It will be interesting to improve the accuracy of our predictions by %using higher resolution, larger extraction radii and especially larger initial %orbital separations. Larger separations, as used in sequence 3, will be necessary to perform accurate comparisons with PN predictions for the evolution of eccentric binaries (see \cite{Hinder2008}). Such an analysis is beyond the scope of this work, however, and a corresponding analysis of the sequence 3 data will be presented elsewhere. For equal masses, we found that eccentric binary mergers with $L\simeq L_{\rm crit}$ maximize the Kerr parameter of the final black hole. Using arguments based on point-particle extrapolations, we proposed two conjectures: (1) $j_{\rm fin}\simeq 0.724\pm 0.13$ should be close to the largest Kerr parameter that can be produced by {\it any} non-spinning black hole binary merger, independently of the binary's mass ratio; and (2) even if we consider maximally spinning holes with spins aligned with the orbital angular momentum, orbital eccentricity should {\it not} lead to violations of the cosmic censorship conjecture. It will be interesting to check these conjectures using numerical simulations of eccentric binaries with unequal masses and non-zero spins. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% | 7 | 10 | 0710.3823 |
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