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1808 | 1808.05170_arXiv.txt | In most astrophysical processes involving synchrotron radiation, the pitch-angle distribution of the electrons is assumed to be isotropic. However, if electrons are accelerated anisotropically, e.g, in a relativistic shock wave with an ordered magnetic field or in magnetic reconnection regions, the electron pitch angles might be anisotropic. In this work, we study synchrotron radiation from electrons with a pitch-angle distribution with respect to a large-scale uniform magnetic field. Assuming that the pitch-angle distribution is normal with a scatter of $\sigma_p$ and that the viewing direction is where the pitch-angle direction peaks, we find that for electrons with a Lorentz factor $\gamma$, the observed flux satisfies $F_\nu\propto\nu^{2/3}$ for $\nu\ll\nu_{\rm cr}$ ($\nu_{\rm cr}$ is the critical frequency of synchrotron), if $\sigma_p\lesssim1/\gamma$ is satisfied. On the other hand, if $\sigma_p\gg1/\gamma$, the spectrum below $\nu_{\rm cr}$ is a broken power law with a break frequency $\nu_{\rm br}\sim2\nu_{\rm cr}/\sigma_p^3\gamma^3$, e.g., $F_\nu\propto\nu^{2/3}$ for $\nu\ll\nu_{\rm br}$ and $F_\nu\propto\nu^{1/3}$ for $\nu_{\rm br}\ll\nu\ll\nu_{\rm cr}$. Thus the ultimate synchrotron line of death is $F_\nu\propto\nu^{2/3}$. We discuss the application of this theory to blazars and gamma-ray bursts (GRBs). | Synchrotron radiation is one of non-thermal radiation mechanisms that plays an important role in many astrophysical sources. The classical synchrotron theory states that the radiation power from an electron is proportional to $\nu^{1/3}$ below a characteristic frequency \citep[e.g.][]{ryb79,jac98}. This feature seems supported by some observations from astrophysical sources. However, we should note that the $\nu^{1/3}$ spectrum from an electron corresponds to the angle-integrated radiation rather than radiation per solid angle along the line of sight. The latter has a spectrum $\nu^{2/3}$ at low frequencies \citep[e.g.][]{ryb79,jac98}. In most astrophysical processes, since an underlying assumption is that the pitch-angle distribution of electrons is isotropic, the classical $\nu^{1/3}$ spectrum is taken as the default. Such a spectrum is regarded as a ``line of death'' for synchrotron radiation \citep[e.g.][]{pre98}, and deviation of this spectrum has been regarded evidence against synchrotron as the radiation mechanism. On the other hand, the momentum distribution of electrons might appear significantly anisotropic if there exists ordered magnetic fields in the emission region. When electrons move helically in a large-scale magnetic field, the pitch-angle distribution would be anisotropic\footnote{The pitch angle here is defined as the angle between the the electron momentum and the magnetic field line.}, which would cause the deviation from the classical $\nu^{1/3}$ spectrum. For the most extreme case that the angle distribution is a Delta function, the observed spectrum would satisfy $F_\nu\propto\nu^{2/3}$ at low frequencies. In this work, we consider synchrotron radiation from electrons with a pitch-angle distribution with respect to a large-scale magnetic field. The theoretical framework is laid out in Section 2. The synchrotron spectra under different conditions are presented in Section 3. Some astrophysical applications are discussed in Section 4, and the results are summarized in Section 5. | In this work, we study synchrotron radiation from electrons with a pitch-angle distribution, and apply it to blazars and GRBs. Following new conclusions are obtained: \begin{itemize} \item { We consider that the sines of pitch angles of the electrons satisfy a normal distribution with a scatter $\sigma_p$, and the line of sight is along the direction where electron distribution is the maximum. Due to an observational selection effect (bright bursts tend to be detected), such a geometry may be the most relevant geometry for observers. We find that for electrons with a Lorentz factor $\gamma$, if $\sigma_p\lesssim1/\gamma$, the observed flux satisfies $F_\nu\propto\nu^{2/3}$ for $\nu\ll\nu_{\rm cr}$; if $\sigma_p\gg1/\gamma$, the spectrum below $\nu_{\rm cr}$ is a broken power law with a break frequency $\nu_{\rm br}\sim2\nu_{\rm cr}/\sigma_p^3\gamma^3$, e.g., $F_\nu\propto\nu^{2/3}$ for $\nu<\nu_{\rm br}$ and $F_\nu\propto\nu^{1/3}$ for $\nu_{\rm br}<\nu<\nu_{\rm cr}$. } \item{ The radiation spectra from astrophysical sources beyond synchrotron death line, i.e., $q>1/3$, can be interpreted by invoking electrons with an extremely anisotropic pith-angle distribution, i.e., $\sigma_p\lesssim1/\gamma_1$, where $\gamma_1$ is the minimum Lorentz factor of the electrons. Such anisotropic pitch angles may be realized by invoking particle acceleration in an ordered magnetic field of a relativistic shock wave or magnetic reconnection region. For example, for a relativistic shock wave with Lorentz factor $\Gamma$, in the shock wave frame, the upstream particles would flow into the shock wave with a Lorentz factor $\Gamma$, and the upstream particles would be confined in an angle of $\lesssim1/\Gamma$. Therefore, if the magnetic field is large-scare uniform (the deflection angle of field line is much less than $\sim1/\gamma_1$), the pitch-angle distribution of particles would be confined within the angle $\lesssim1/\Gamma$. } \item{The theory may be applied to blazars and GRBs. For radio emission of blazars, since the minimum Lorentz factor of electrons could be small, e.g., $\gamma\sim1-10$, the requirements for the pitch-angle distribution and the deflection of magnetic field lines are modest. It is therefore relatively easy to break the synchrotron death line, and the theory can interpret some blazars whose low-energy spectral index is harder than $q=1/3$. For prompt gamma-ray emission of GRBs, since the minimum Lorentz factor of electrons is large, e.g., $\gamma\sim10^4$, the conditions under which our theory applies is more demanding. In any case, if those conditions are satisfied, GRBs that are slightly beyond the traditional synchrotron line of death can be in principle still accounted for within the framework of synchrotron radiation model (e.g. within the ICMART model). The ultimate line of death is pushed to $q=2/3$. } \end{itemize} | 18 | 8 | 1808.05170 |
1808 | 1808.09168_arXiv.txt | { The KATRIN experiment will determine the effective electron anti-neutrino mass with a sensitivity of $200\,\mathrm{meV/c}^2$ at 90\% CL. The energy analysis of tritium $\beta$-decay electrons will be performed by a tandem setup of electrostatic retarding spectrometers which have to be operated at very low background levels of $<10^{-2}\,$counts per second. This benchmark rate can be exceeded by background processes resulting from the emanation of single ${}^{219,220}$Rn atoms from the inner spectrometer surface and an array of non-evaporable getter strips used as main vacuum pump. Here we report on a the impact of a cryogenic technique to reduce this radon-induced background in electrostatic spectrometers. It is based on installing a liquid nitrogen cooled copper baffle in the spectrometer pump port to block the direct line of sight between the getter pump, which is the main source of ${}^{219}$Rn, and the sensitive flux tube volume. This cold surface traps a large fraction of emanated radon atoms in a region outside of the active flux tube, preventing background there. We outline important baffle design criteria to maximize the efficiency for the adsorption of radon atoms, describe the baffle implemented at the KATIRN Pre-Spectrometer test set-up, and report on its initial performance in suppressing radon-induced background. } | \label{sec:instruction} The KATRIN (\textbf{KA}rlsruhe \textbf{TRI}tium \textbf{N}eutrino) experiment \cite{KATRINDesignReport} is currently being commissioned at the Karlsruhe Institute of Technology (KIT) in Germany. The experiment will investigate the kinematics of the tritium $\beta$-decay close to the kinematic endpoint of $E_{0} \approx 18.6\,$keV in a model-independent way, relying on energy and momentum conservation only. It will determine the effective electron anti-neutrino mass $m_{\nu}$ with a sensitivity of $200\,\mathrm{meV/c}^2$ at 90\% confidence level after three live years of measurements (equivalent to five calender years) \cite{drexlin2013current}. The imprint of a non-zero neutrino mass is a distortion of the shape of the electron energy spectrum a few eV below the endpoint. Due to the rather low signal countrate of a few times $10^{-2}\,$counts per second (cps) in this narrow region-of-interest, a low background rate of $\leq10^{-2}\,$cps has to be achieved to obtain a good signal to background ratio. A schematic overview of the $70\,\mathrm{m}$ long KATRIN setup is given in Fig.~\ref{fig:beamline}. A windowless gaseous tritium source (WGTS) \cite{WGTS} will provide an unprecedented activity of $\sim 10^{11}$ $\beta$-decays per second. Only electrons emitted in the WGTS into the forward cone are guided adiabatically by strong magnetic fields ($\sim 5\,$T) through the transport section to the spectrometers. The electrons are constrainted to their magnetic flux tube throughout the experiment. The two transport elements, the differential pumping section (DPS) \cite{Kosmider:2012:PHD} and the cryogenic pumping section (CPS) \cite{CPS}, will remove residual tritium gas by turbomolecular pumps (DPS) and cryotraps (CPS). The spectrometer section, consisting of a Pre- and a larger Main Spectrometer, analyzes the electron energy. The Pre-Spectrometer offers the option to act as a filter, where only electrons with energies close to $E_0$ would be transmitted to the Main Spectrometer for precision energy filtering, thereby limiting the incoming flux of electrons. Both spectrometers are operated as MAC-E filters (\cite{Bea80}, \cite{Lobashev1985}, and \cite{picard92}). \begin{figure}[t] \centering \includegraphics[width=\textwidth]{figures/fig1_beamline.jpg} \caption[KATRIN setup]{Schematic overview of the $70\,\mathrm{m}$ long KATRIN setup: \textbf{a} rear section, \textbf{b} windowless gaseous molecular tritium source (WGTS), \textbf{c} differential pumping section (DPS), \textbf{d} cryogenic pumping section (CPS), \textbf{e} Pre-Spectrometer, \textbf{f} Main Spectrometer, \textbf{g} detector system.} \label{fig:beamline} \end{figure} In this filter technology, the magnetic field strength in the spectrometer drops by several orders of magnitude between either end and the center of the vessel (see Fig.~\ref{fig:setup}). In the Main Spectrometer, this field ratio amounts to $1/20000$. Due to the conservation of the magnetic moment of the electrons and the adiabatic field layout, the magnetic gradient force transforms almost the entire cyclotron energy of the isotropically emitted electrons into longitudinal energy, which can then be analyzed by the electrostatic retarding potential. This \textbf{M}agnetic \textbf{A}diabatic \textbf{C}ollimation with an \textbf{E}lectrostatic filter, the MAC-E principle, allows the spectrometers to act as a high-pass filter for electrons. Accordingly, only those electrons that are able to overcome the potential barrier in the spectrometer section are counted by a monolithic segmented Si-PIN diode array \cite{Amsbaugh2015} at the downstream end of the Main Spectrometer. By measuring at different filter voltages, the shape of the integrated energy spectrum of tritium $\beta$-decay in the vicinity of $E_{0}$ can be determined (for more details, see \cite{drexlin2013current}). In order to reach the required background rate of $10^{-2}\,$cps, low-energy secondary electrons from radioactivity and cosmic-ray muon interactions in the vessel walls have to be shielded by magnetic and electrostatic fields that exhibit a high degree of axial symmetry \cite{KATRINDesignReport}. However, this technique does not shield against electrons from radioactive decays of neutral unstable atoms which are emanated from the inner spectrometer walls or from auxiliary equipment, such as pumps required to maintain the stringent ultra-high vacuum (UHV) conditions of $p \sim 10^{-11}\,$mbar. \begin{figure} [t] \centering \includegraphics[width=\textwidth]{figures/fig2_macefilter3.pdf} \caption[the pre-spectrometer MAC-E filter]{ Electromagnetic layout of the Pre-Spectrometer with two superconducting magnets providing the characteristic flux tube of the magnetic field (blue lines) that guides electrons (red line) from the source to the detector through the electro-static field of the MAC-E filter. This configuration is also responsible for the inherent magnetic bottle characteristics of a MAC-E filter trapping charged particles (green line) which are produced in the low-field region. A concentric inner electrode system of thin stainless steel cones (black line) and wires (broken lines) allows to fine-tune the retarding potential.} \label{fig:setup} \end{figure} Previous investigations \cite{RadonPaper, NuclearDecayPaper} have revealed that the $\alpha$-decay of ${}^{219,220}$Rn atoms, emanated in small quantities from the installed non-evaporable getter (NEG) pump and from the inner surface of the stainless steel vessel, constitute a major source of background in KATRIN which cannot be shielded by electromagnetic fields. This background arises from electrons with energies from a few eV up to a few hundred keV \cite{BROWNE2001763, nancy} which accompany the primary $\alpha$-decay of both radon isotopes. Subsequent investigations focus on actively removing trapped high-energy electrons \cite{ECRpaper} and on the signature of this background by detailed simulations \cite{nancy2} and specific measurements \cite{RadonPaper}. Here we report on a passive background reduction technique that traps emanated radon atoms on a large cold surface located outside the sensitive volume of the magnetic flux tube. By confining the neutral radon atoms to the surface of the cold trap for an extended period of time, electrons produced by the subsequent $\alpha$-decay of the absorbed radon atom are magnetically shielded and cannot contribute to background processes. Capturing radon atoms on charcoal-covered cryogenic panels is a commonly used method for low count rate experiments to reduce the radon-induced background rate. This method works well even for ${}^{222}$Rn with a half-life of 3.8\,days. However, in tests we found that the fine dust of charcoal released from cryo-panels can be deposited in the vicinity of the panels. With large numbers of tiny ceramic insulators inside the KATRIN spectrometers, this can lead to unwanted leakage currents. Since we are mainly concerned about the short-lived isotopes ${}^{219,220}$Rn, our novel design employs cryo-panels made of pure OFHC copper. The mean sojourn time of a radon atom captured on the cold surface has to be long enough to allow the radon isotope to decay while still attached to the cryo-panel. The paper is organized as follows: First we describe the Pre-Spectrometer in its stand-alone test setup configuration (section~\ref{sec:Pre-Spectrometer}). Next we outline the key design criteria for a baffle to optimize the efficiency in adsorbing emanated radon atoms, and depict the configuration as implemented in one of the pump ports of the Pre-Spectrometer (section~\ref{sec:baffle}). We then report on different measurements with this improved setup, which demonstrate the feasibility of our method by significantly reducing the observed rate of radon-induced background (section~\ref{sec:baffle_measurement}). Finally, we give an outlook on how to implement this cryogenic technology at the much larger Main Spectrometer (section~\ref{sec:outlook}). | \label{sec:outlook} The short-lived radon isotopes ${}^{219}$Rn and ${}^{220}$Rn which emanate from the porous NEG strips and the surface of the spectrometer are a serious source of background in electrostatic retarding spectrometers due to the emission of electrons with energies up to 200 keV. As a result of the inherent magnetic trap formed by the MAC-E-filter, single $\alpha$-decay processes of these isotopes result in enhanced levels of background over an extended period of time, adding a non-Poissonian background component \cite{NuclearDecayPaper} to the cosmic-ray induced background. Thus, radon emanation seriously limits the neutrino mass sensitivity of KATRIN if no countermeasures are taken. Due to the fact that the emanating radon atoms are neutral, the highly efficient magnetic and electrostatic shielding against charged particles from the vessel walls is ineffective. This deficiency calls for additional active and passive background suppression techniques. \begin{figure}[t] \centering \includegraphics[width=9.5cm]{figures/fig11_ringcount.pdf} \caption[ring counter]{The number of detector hits as a function of the event duration for the three measurement campaigns (warm baffle = red, cold baffle = blue, blank run = green) can be used to separate internal conversion electrons from ${}^{219}$Rn $\alpha$-decay from radon spikes initiated by electrons from shake-off and Auger processes following $\alpha$-decays of both ${}^{219}$Rn and ${}^{220}$Rn. For each background class, consistent suppression factors provided by a cold baffle are obtained.} \label{fig:ringcount} \end{figure} Here we have reported on the implementation of a passive method, in form of a LN$_2$-cooled baffle that traps radon emanated from the NEG pump and the inner vessel walls. By comparing measurements with a LN$_2$-cooled baffle to measurements with a warm baffle, an overall singles background reduction factor of $(5.4 \pm 2.7)$ has been observed. By identifying individual radon spikes from high-energy trapped electrons, a comparable background reduction factor of $(4.4 \pm 1.2)$ was observed, giving a first and consistent proof-of-principle of this novel background reduction technique in MAC-E filters. The Main Spectrometer ($10^{-11}\,\mathrm{mbar}$) was originally equipped with a scaled-up version of the NEG pump described above, with $3\,$km of St707 getter strips installed. By extrapolating the results of the PS tests reported here and in \cite{nancy2}, the expected radon-induced background rate due to the NEG material and the large surface of the Main Spectrometer was estimated to be of the order of $1\,\mathrm{cps}$ \cite{Fraenkle:2010:PHD, NuclearDecayPaper, Mertens:2012:PHD}. This value was confirmed by the commissioning measurements (\cite{Goerhardt:2014:PHD}, \cite{LRT2017}). As such a large background rate would significantly reduce the sensitivity of KATRIN on $m_{\nu}$ \cite{NuclearDecayPaper, Mertens:2012:PHD, Wandkowsky:2013:PHD}, this background class has to be further reduced by more than two orders of magnitude. Based on the encouraging results of this paper, a large-scale baffle system to reduce radon-induced background in the Main Spectrometer has been designed and optimized along the design criteria outlined above. In order to reach a similar pumping speed for radon in the much larger Main Spectrometer (690\,m$^2$), the cold surface would have to be 28 times larger (required: 5.5\,m$^2$ cold surface area, actual design: 6.6\,m$^2$). In addition, one has to take into account the 7 times longer time of flight between two hits of the wall, which would increase the decays in flight by approximately the same factor. For radon-induced background rates in the PS of up to $2\cdot 10^{-3}$\,cps, the scaled-up rate in the Main Spectrometer is still close enough to the limit proposed in the KATRIN Design Report \cite{KATRINDesignReport}. For the Main Spectrometer a modified design was adopted for the baffles. It retains a sufficiently large effective pumping speed for hydrogen and tritium (40\% of the value without baffles) by abandoning the central disk in the PS baffle design \cite{lit:SDSvacuum16}. The large Main Spectrometer baffles were made OFHC copper \cite{RadonAdsorption}. In principle, the baffle operation and radon adsorption times $\tau_{\mathrm{des}}$ could be further improved by gold-plating the V-shaped copper panels or by annealing the base material (copper) at a temperature of $\sim 1000\,$K \cite{RadonAdsorption}. A detailed description of this large-scale baffle and its performance is given in \cite{Goerhardt:2014:PHD, MSradonsimu2017}. In addition to the passive background reduction via the technique of a liquid-nitrogen-cooled copper baffle discussed here, further active suppression methods of background induced by magnetically stored electrons were investigated, such as a cyclic application of an electric dipole \cite{Hilk:2017:Thesis} in combination with a magnetic pulse \cite{Wandkowsky:2013:PHD, Behrens:2017:Thesis, arXivBehrens2018}, and electron cyclotron resonance \cite{ECRpaper, Mertens:2012:PHD}. The method of removing emanated Rn atoms from the active flux tube of a MAC-E filter is studied in more detail in measurements with the Main Spectrometer \cite{Goerhardt:2014:PHD, LRT2017}. By combining all the background reduction techniques, a staged approach to reach conditions where the spectrometer is essentially free of radon-induced background is at hand for the KATRIN neutrino mass measurements. | 18 | 8 | 1808.09168 |
1808 | 1808.00080_arXiv.txt | Binary pulsars allow us to carry out precision tests of gravity and have placed stringent bounds on a broad class of theories beyond general relativity. Current and future radio telescopes, such as FAST, SKA, and MeerKAT, may find a new astrophysical system, a pulsar orbiting around a black hole, which will provide us a new source for probing gravity. In this paper, we systematically study the prospects of testing general relativity with such black hole-pulsar binaries. We begin by finding a mapping between generic non-Einsteinian parameters in the orbital decay rate and theoretical constants in various modified theories of gravity and then summarize this mapping with a ready-to-use list. Theories we study here include scalar-tensor theories, varying $G$ theories, massive gravity theories, generic screening gravity and quadratic curvature-corrected theories. We next use simulated measurement accuracy of the orbital decay rate for black hole-pulsar binaries with FAST/SKA and derive projected upper bounds on the above generic non-Einsteinian parameters. We find that such bounds from black hole-pulsars can be stronger than those from neutron star-pulsar and neutron star-white dwarf binaries by a few orders of magnitude when the correction enters at negative post-Newtonian orders. By mapping such bounds on generic parameters to those on various modified theories of gravity, we find that one can constrain the amount of time variation in Newton's constant $G$ to be comparable to or slightly weaker than than the current strongest bound from solar system experiments, though the former bounds are complementary to the latter since they probe different regime of gravity. We also study how well one can probe quadratic gravity from black hole quadrupole moment measurements of black hole-pulsars. We find that bounds on the parity-violating sector of quadratic gravity can be stronger than current bounds by six orders of magnitude. These results suggest that a new discovery of black hole-pulsars in the future will provide powerful ways to probe gravity further. | \label{sec:intro} General relativity (GR) is currently the most well-tested theory of gravity. Nevertheless, due to its inconsistency with quantum mechanics, it is only an effective field theory below some energy threshold. Eventually, at some length or energy scales, one expects to find deviations from GR predictions. So far, GR has passed all the tests with flying colors~\cite{Will:2014kxa}. Thus, it is important to probe even deeper to search for non-GR effects using new sources that we hope to detect with current and future detectors. Probing GR can also shed light on cosmology as another motivation to consider going beyond GR is to explain current accelerating expansion of our universe and missing mass problem in galaxies without introducing dark energy or dark matter~\cite{Jain:2010ka,Clifton:2011jh,Joyce:2014kja,Koyama:2015vza}. Various probes of gravity can be classified into weak-field and strong-field tests~\cite{Psaltis:2008bb,Baker:2014zba,Yunes:2016jcc}. Solar system experiments~\cite{Will:2014kxa} and cosmological observations~\cite{Jain:2010ka,Clifton:2011jh,Joyce:2014kja,Koyama:2015vza} fall into the weak-field category, as the amount of curvature and gravitational potential that these experiments and observations probe is weak and small. On the other hand, recent direct gravitational-wave (GW) measurements from binary black hole (BH) mergers~\cite{Abbott:2016blz,Abbott:2016nmj,TheLIGOScientific:2016pea,Abbott:2017vtc,Abbott:2017oio,Abbott:2017gyy} allow us to probe the strong and dynamical field regime of gravity for the first time~\cite{TheLIGOScientific:2016src,Yunes:2016jcc}. Precision tests of gravity have also been carried out via binary pulsar (PSR)\footnote{In this paper, binary PSR refers to a binary consisting of a PSR with a \emph{stellar} companion.} observations~\cite{stairs,WexRadioPulsars} (see e.g.~\cite{Archibald:2018oxs} for a very recent work on testing the strong equivalence principle with the triple system). Such systems allow us to probe both the weak-field and strong-field effects. This is because these binaries are widely-separated relative to GW sources of compact binaries that are about to coalesce (and hence weak-field), and at the same time, PSRs are neutron stars which are very compact objects (and thus strong-field). Such systems (PSRs with white dwarf (WD) companions in particular) are ideal for probing scalar dipole emission in scalar-tensor theories~\cite{Freire:2012mg} that is absent in GR. One can use parameterized post-Keplerian (PPK) parameters, such as the advance rate of periastron, orbital decay rate and Shapiro time delay, to carry out precision tests of GR. Measurements of any two PPK parameters determine the masses of compact objects in binary PSRs, while any additional PPK parameter measurements further probe the consistency of gravitational theory. Binary BHs and binary neutron stars have been found via electromagnetic-wave~\cite{stairs,Valtonen:2008tx,WexRadioPulsars} and GW observations~\cite{Abbott:2016blz,Abbott:2016nmj,TheLIGOScientific:2016pea,Abbott:2017vtc,Abbott:2017oio,Abbott:2017gyy,TheLIGOScientific:2017qsa}. Interestingly, a binary that consists of one BH and one neutron star has not been found yet. One possibility to look for such systems is through GW observations, which are useful for probing non-GR theories such as scalar-tensor theories~\cite{berti-buonanno,Yagi:2009zm,Yagi:2009zz,Zhang:2017sym}. Radio observations offer another possible way of finding a PSR orbiting around a BH. Such a system may be found with either the Five-hundred-meter Aperture Spherical radio Telescope (FAST) that is undergoing commissioning, MeerKAT or a next-generation radio telescopes such as the Square Kilometer Array (SKA). Population synthesis suggests that there may be around 3--80 BH-PSRs in the Galactic disk and FAST may detect up to 10\% of them~\cite{Shao:2018qpt}. Radio telescopes may detect two different types of PSR in the BH-PSR binary. A normal PSR is a younger, slower spinning PSR. A millisecond PSR (MSP) is older and recycled. MSPs are ideal for testing GR because they spin significantly faster and achieve a better timing precision than normal PSRs (100ns vs 100$\mu$s respectively) \cite{QuadrupoleMeasureabilitySgrA_PsaltisWex}. This better timing precision allows for a more accurate system measurement. Note that MSPs are usually harder to find than normal pulsars due to selection effects. This is particularly true for the Galactic center, where the large distance and scattering effects in the interstellar medium may have prevented a discovery so far. Regarding formation of a BH-MSP binary, at least two possible scenarios exist. First, an exchange interaction of a binary can create a BH-MSP binary inside a dense region such as a globular cluster or the Galactic center. Second, a BH-MSP binary can evolve directly from a finely tuned initial system of main sequence stars. For example, when the masses of initial stars are comparable, a neutron star can form first, which is being accreted by a companion and spun up, and eventually the companion collapses to a BH. Thus, while binaries with a BH and normal PSR are more likely, BH-MSP binaries may still be found. BH-PSR systems are powerful for testing GR, including no-hair properties of BHs~\cite{QuadrupoleBoundFunctionofPb,WexBHP,QuadrupoleMeasureabilitySgrA_PsaltisWex}, scalar-tensor theories~\cite{WexBHP,Wex_ScalarTensor}, higher-curvature theories~\cite{EDGBBoundComparison}, higher-dimensional gravity~\cite{Simonetti:2010mk} and quantum gravity effects~\cite{Estes:2016wgv}. The reason is as follows. The relative velocity of a binary is given by $v = (2 \pi M / P)^{1/3}$ with $c=G=1$, where $M$ is the total mass while $P$ is the orbital period. For a PSR orbiting around a stellar-mass BH, $M$ is larger than that of PSR binaries with NS or WD companions, but one expects $P$ to be also larger. In fact, the BH-PSR relative velocity is smaller than the NS-PSR or PSR-WD case (typically by a factor of 2) because the longer period more than compensates for the larger total mass. Additionally, the measurement accuracy of the orbital decay rate is expected to be similar to that of NS-PSR or PSR-WD. Thus, BH-PSR systems will have more advantage on probing non-GR effects that enter at a negative post-Newton (PN) order\footnote{PN expansion assumes that the orbital motion is sufficiently slow relative to the speed of light. A correction term is said to be of relative $n$PN order when it is proportional to $(v/c)^{2n}$ relative to the leading contribution in GR.}, such as scalar dipole radiation in scalar-tensor theories. In this paper, we study in more detail how well one can probe modifications to GR in theories that have not yet been studied in the context of BH-PSRs. The first half of the paper focuses on using the orbital decay rate measurement. We will first introduce a generic parameterization that captures the non-GR modifications to the orbital decay rate~\cite{Yunes:2010qb}. We will next derive projected bounds on this parameter from a BH-MSP with FAST and SKA based on a simulated measurement accuracy in~\cite{WexBHP}. (This simulation will be discussed further in Sec.~\ref{sec:Pdot-general}.) We will then compare such bounds to those on parameterized post-Einsteinian (PPE)~\cite{Yunes:2009ke} parameters, which capture the non-GR modifications in gravitational waveforms from compact binary mergers. The PPE formalism (or its modified version) has already been applied to recent GW events~\cite{TheLIGOScientific:2016src,KentGWbounds}. We will further map such generic bounds to those on specific modified theories of gravity by creating a ``dictionary'' between the generic and theoretical parameters shown in Table~\ref{tab:GammaValues}. In particular, we will study theories with time varying gravitational constant~\cite{Will:2014kxa}, Lorentz-violating graviton mass~\cite{Finn:2001qi}, Lorentz-preserving graviton mass~\cite{MassiveGravityCG}, and generic screening mechanisms~\cite{Zhang:2018dxi}. { \newcommand{\minitab}[2][l]{\begin{tabular}{#1}#2\end{tabular}} \renewcommand{\arraystretch}{1.6} \begingroup \squeezetable \begin{table*}[htb] \begin{centering} \begin{tabular}{c|c|c|c|c|c|c} \hline \hline \multirow{2}{*}{Theory}&\multirow{2}{*}{$\gamma$}&\multirow{2}{*}{$f(e)$}&\multirow{2}{*}{$n$}&Theoretical&\multirow{2}{*}{Refs.}&Stronger\\ & & & & parameters & & bounds? \\ \hline \hline \multirow{2}{*}{Time-Varying G (Sec.~\ref{subsec:VaryingG})}&\multirow{2}{*}{ $\frac{5 \dot G M^3}{48 m_{c} m_{p}} \left[1-s_{p}^{\dot G} \left(1+\frac{m_{c}}{M}\right)-s_{c}^{\dot G} \left(1+\frac{m_{p}}{M}\right)\right]f(e)$ }&\multirow{2}{*}{$\frac{1}{F_\GR(e)}$}& \multirow{2}{*}{$-4$} &\multirow{2}{*}{$\dot G / G$ }&\multirow{2}{*}{\cite{WexRadioPulsars}} &\multirow{2}{*}{\checkmark} \\ & & & & && \\ \hline Lorentz-violating & \multirow{2}{*}{ $\frac{5M^2}{24 } m_g ^2 f(e)$}&\multirow{2}{*}{ $\frac{1}{(1-e^2)^{1/2}F_\GR(e)}$}&\multirow{2}{*}{$-3$}&\multirow{2}{*}{$m_g$}& \multirow{2}{*}{\cite{Finn:2001qi}}& \multirow{2}{*}{\xmark} \\ Massive Gravity (Sec.~\ref{subsec:FierzPauli}) & & & & && \\ \hline Cubic Galileon & \multirow{2}{*}{$\frac{25}{32} \pi \lambda ^2 \frac{ M_\PL M_\Q^2 M^{3}}{m_p^2 m_c^2} m_g f(e)$}&\multirow{2}{*}{$\frac{F_\CG(e)}{F_{\GR}(e)}$}&\multirow{2}{*}{$-11/4$}&\multirow{2}{*}{$m_g$}&\multirow{2}{*}{\cite{MassiveGravityCG}}& \multirow{2}{*}{\xmark} \\ Massive Gravity (Sec.~\ref{subsec:CubicGalileon}) & & & & & & \\ \hline General Screen Modified & \multirow{2}{*}{$\frac{5}{192} (\epsilon_p-\epsilon _c)^2 f(e)$} & \multirow{2}{*}{$\frac{F_\SMG} {F_\GR(e)}$}& \multirow{2}{*}{$-1$}& \multirow{2}{*}{$\phi_\VEV / M_\PL$}& \multirow{2}{*}{\cite{Zhang:2018dxi}}& \multirow{2}{*}{\xmark} \\ Gravity (Sec.~\ref{subsec:SMG}) & & & & & & \\ \hline \multirow{2}{*}{(massless) Scalar-Tensor} & \multirow{2}{*}{$ \frac{5}{96} \left(\bar\alpha_p^\ST-\bar\alpha_c^\ST\right)^2 f(e)$} & \multirow{2}{*}{$\frac{F_\SMG}{F_\GR(e)}$}& \multirow{2}{*}{$-1$} & \multirow{2}{*}{($\alpha_0$, $\beta_0$)}& \multirow{2}{*}{\cite{Freire:2012mg}} & \multirow{1}{*}{\checkmark} \\ & & & & & & \cite{WexBHP,Wex_ScalarTensor} \\ \hline \multirow{2}{*}{Einstein-dilaton Gauss-Bonnet} & \multirow{2}{*}{$\frac{5 \pi}{24} \left(\bar\alpha_p^\EdGB-\bar\alpha_c^\EdGB\right)^2f(e)$}&\multirow{2}{*}{---}&\multirow{2}{*}{$-1$} &\multirow{2}{*}{$\sqrt{\alpha_\EdGB}$}&\multirow{2}{*}{\cite{EDGBBoundComparison}} & \multirow{1}{*}{\checkmark} \\ & & & & & & \cite{EDGBBoundComparison} \\ \hline \multirow{2}{*}{Einstein-\AE ther} & \multirow{2}{*}{$ \frac{5 \mathcal{C}_\EA }{32 } \left(1-\frac{c_{14}}{2}\right) \left(s_{p}^\EA-s_{c}^\EA \right)^2f(e)$} & \multirow{2}{*}{---} &\multirow{2}{*}{$-1$}&\multirow{2}{*}{($c_{+}$, $c_{-})$}&\multirow{2}{*}{\cite{Yagi:2013ava}} & \multirow{2}{*}{?} \\ & & & & & & \\ \hline \multirow{2}{*}{Khronometric} & \multirow{2}{*}{$ \frac{5 \mathcal{C}_\KG}{32 } \left(1-\frac{\alpha_\KG}{2}\right) \left(s_{p}^\KG-s_{c}^\KG\right)^2 f(e)$} & \multirow{2}{*}{---} & \multirow{2}{*}{$-1$}& \multirow{2}{*}{($\lambda_\KG$, $\alpha_\KG$, $\beta_\KG$)}& \multirow{2}{*}{\cite{Yagi:2013ava}}& \multirow{2}{*}{?} \\ & & & && & \\ \hline \hline \end{tabular} \end{centering} \caption{ \label{tab:GammaValues} Mapping between non-GR parameters ($\gamma$ and $n$) in the orbital decay rate $\dot P$ in Eq.~\eqref{eq:generic-Pdot} to theoretical parameters in various example modified theories of gravity, together with some references. These expressions are valid for any compact binaries (not specific to BH-PSRs). The first four theories are those considered in Sec.~\ref{sec:Pdot-example}, while the last four theories are presented only for reference. Note that we study bounding EdGB gravity in Sec.~\ref{subsec:EdGB} via BH quadrupole moment measurement which is different from the orbital decay rate presented here. Theoretical parameters are presented in the fifth column. The last column shows whether BH-PSR bounds are stronger than other existing bounds (\checkmark: yes; \xmark: no; ?: unknown). The meaning of each parameter in the second column is as follows. $m_p$: primary PSR's mass, $m_c$ companion's mass, $M$: total system mass, $e$: eccentricity, $M_\PL$: Planck mass, $M_\Q$: a mass parameter in Eq.~\eqref{eq:MQdef}, $\lambda$: a numerical constant in Eq.~\eqref{eq:lambdadef}, $\epsilon_A$: a screening parameter in SMG in Eq.~\eqref{eq:ScreenScalarCharge}, $C_\EA$ and $C_\KG$: a function of theory parameters in Eqs.~(114) and (124) of~\cite{Yagi:2013ava} for Einstein-\ae ther and khronometric theories respectively, $c_{14}$ and $\alpha_\KG$: a combination of coupling constants in Einstein-\ae ther theory and khronometric theory respectively. $\bar\alpha_A$ is the scalar charge. In many of scalar-tensor theories, it is non-vanishing for stars while it is zero for a BH~\cite{Hawking:1972qk,Bekenstein:1995un,Sotiriou:2011dz}. In EdGB gravity, such a charge vanishes for stars~\cite{Yagi:2011xp,EDGBBoundComparison} while that for a BH is in Eq.~(37) of~\cite{Berti:2018cxi}. $s_A$ is the sensitivity and that for a neutron star in Einstein-\ae ther and khronometric theory has been computed in~\cite{Yagi:2013qpa,Yagi:2013ava}\footnote{The fitting function for the NS sensitivity in Einstein-\ae ther and khronometric theory can be found in Eq.~(186) or~(C1) of~\cite{Yagi:2013ava}, though the parameter region in which the fit is valid has mostly been ruled out by GW170817~\cite{Gumrukcuoglu:2017ijh,Oost:2018tcv}. } while that for a BH has not been calculated yet. The eccentricity dependent function $f(e)$ is presented in the third column if known, while ``---'' means that the correction has been calculated only for circular binaries ($f=1$). $F_\GR$ is the eccentricity dependence in GR in Eq.~\eqref{eq:FGR}, while $F_\CG(e)$ and $F_\SMG(e)$ are that in cubic Galileon massive gravity and generic screened massive gravity defined in Eqs.~\eqref{eq:FCGdef} and~\eqref{eq:F_SMG} respectively. } \end{table*} \endgroup The second half of this paper focuses on using the BH quadrupole moment measurement with BH-PSRs to probe gravity. A non-vanishing quadrupole moment causes a periodic variation in the PSR motion~\cite{Wex:1998wt}, which can be extracted from the Roemer time delay measurement. For stellar-mass BH-MSP systems, one may be able to measure the BH quadrupole moment within ~10\% accuracy~\cite{WexBHP}, and the accuracy may be 10\% if one finds a PSR orbiting around Sgr~A$^*$~\cite{QuadrupoleBoundFunctionofPb,QuadrupoleMeasureabilitySgrA_PsaltisWex}. (The simulations used for these are discussed in~Sec.~\ref{sec:Q-measurement}.) We apply these projected measurements to quadratic curvature theories, namely Einstein-dilaton Gauss-Bonnet (EdGB) gravity~\cite{Metsaev:1987zx,Maeda:2009uy} and dynamical Chern-Simons (dCS) gravity~\cite{Jackiw:2003pm,Alexander:2009tp} for the even-parity and odd-parity sector respectively. Both theories are motivated from string theory. Analytic BH solutions with arbitrary spin in these theories have not been found yet. Non-rotating and slowly-rotating analytic BH solutions have been constructed in~\cite{Mignemi:1992pm,Yunes:2011we,Sotiriou:2014pfa,EDGBBoundSpin2,EDGBBoundSpin4} for EdGB and in~\cite{Yunes:2009hc,Konno:2009kg,DCSBoundSpin2,DCSBoundspin4} for dCS. \subsection{Executive Summary} \begin{figure}[t] \begin{center} \includegraphics[width=\linewidth]{Gamma_Ratio.pdf} \caption{\label{fig:gammaratio} The ratio in the upper bound on the fractional non-GR correction $\gamma$ to the orbital decay rate between BH-PSR and double PSR systems, as a function of at which PN order the correction enters. This bound uses a millisecond PSR-BH binary with simulations from \cite{WexBHP}. This figure shows how much improvement one finds by using the BH-PSR system compared to the double PSR one in terms of testing GR. For example, if the ratio is below unity (horizontal magenta dotted dashed line), the bound from the former is stronger than that from the latter. The ratio is shown for FAST (red dotted) and SKA (black dotted). Notice that FAST (SKA) has a bound that is an order of magnitude stronger at $-4$ ($-3.5$) PN than the double PSR. At $-4$PN order which corresponds to corrections due to e.g.~time variation in the gravitational constant $G$, the BH-PSR bounds are stronger than the double PSR one by almost two orders of magnitude. } \end{center} \end{figure} We now give a brief summary of this paper. In examining the prospects of bounding non-GR theories in a BH-PSR binary, it is important to compare to the existing method of binary PSR measurements. Figure~\ref{fig:gammaratio} presents the upper bound on the fractional non-GR correction $\gamma$ to the orbital decay rate with BH-PSRs relative to those with the double PSR, as a function of the entering PN order correction. If the ratio is below unity, BH-PSR bounds are stronger than the double PSR ones. Observe that the former can be stronger by orders of magnitude than the latter for negative PN corrections. Next, we apply such bounds on a generic non-GR parameter for the orbital decay rate to specific non-GR theories based on Table~\ref{tab:GammaValues}. For example, Fig.~\ref{fig:VaryingG} presents the upper bound on the time variation in $G$ as a function of the orbital period of BH-PSRs. Observe that BH-PSR observations with SKA are slightly weaker than the current strongest bounds from solar system experiments of NASA Messenger. However, since binary PSRs are sensitive to self gravity effects in the strong gravity regime, their tests complement a weak field test of $\dot G$ such as NASA Messenger. In the strong gravity regime, the time variation in $G$ can be magnified by a factor of 20 from that of weak field tests due to effects of the object's sensitivities \cite{WexBHP}. Binary PSR bounds also provide an independent test for non-GR effects. For other theories that we study in this paper, we find that BH-PSR $\dot P$ bounds are weaker than those obtained from NS-PSR or PSR-WD observations. \begin{figure}[!h] \begin{center} \includegraphics[width=\linewidth]{Varying_Gravitational_Constant.pdf} \caption{\label{fig:VaryingG} Projected bounds on the time variation of the gravitational constant $G$ over the static gravitational constant as a function of the orbital period. The bound is shown for BH-PSR systems with FAST (red dotted) and SKA (black dashed). We also present two solar system bounds with The NASA Messenger (purple dotted-dashed line) \cite{2018NatCo...9..289G} and the Mars ephemeris (blue dotted-dashed line)~\cite{Will:2014kxa}, and the current strongest binary pulsar bound (dot dash orange) \cite{Zhu:2018etc,Zhu:2015mdo}. Notice that SKA can produce bounds that are comparable to or weaker than solar system ones, though the former are complementary to the latter as the two bounds probe different regime of gravity. } \end{center} \end{figure} Regarding bounds on quadratic gravity via BH quadrupole moment measurements, we find that bounds on dCS gravity can be improved by six to seven orders of magnitude for stellar-mass BH-PSR binaries. We also investigate bounding dCS gravity with a PSR orbiting Sgr~A$^*$, but we show that this system cannot reach a tight enough bound to satisfy the small coupling approximation. On the other hand, BH-PSR bounds on EdGB gravity are weaker than the current bounds from e.g.~BH-low-mass X-ray binaries (LMXB)~\cite{Yagi:2012gp} by an order of magnitude. The rest of the paper is organized as follows. In Sec.~\ref{sec:Pdot}, we focus on orbital decay rate measurements. After discussing a generic formalism for describing non-GR corrections to the orbital decay rate and its relation to the PPE formalism in Sec.~\ref{sec:Pdot-general}, we study BH-PSR bounds in various modified theories of gravity in Sec.~\ref{sec:Pdot-example}. In Sec.~\ref{sec:QuadrupoleMoment}, we focus on BH quadrupole moment measurements. After reviewing such measurements for BH-PSRs in Sec.~\ref{sec:Q-measurement}, we study bounds on two kinds of quadratic gravity in Sec.~\ref{sec:Q-example}. We conclude in Sec.~\ref{sec:conclusion} and give possible avenues for future work. We use the geometric units of $c=1$ and $G=1$ throughout the paper unless otherwise stated. | \label{sec:conclusion} In this paper, we studied how well one can probe gravity with orbital decay rate and BH quadrupole moment measurements using future BH-PSR observations. Regarding the former, we showed that bounds from generic non-GR modifications to $\dot P$ can be stronger than those from the double PSR by a few orders of magnitude, especially for corrections entering at negative PN orders. We mapped this result to various example modified theories of gravity and found that e.g.~bounds on $\dot G$ can be much stronger than the current bound from solar system experiments. Regarding the latter, we showed that the Roemer time delay for certain stellar-mass BH-PSR configurations can be used to place bounds on dCS gravity that are six orders of magnitude stronger than the current most stringent bounds. Thus, the detection of a BH-PSR binary will allow new tests of gravity. Future work includes extending BH-PSR bounds to Lorentz-violating theories, such as Einstein-\ae ther~\cite{Jacobson:2000xp,Jacobson:2008aj} and khronometric~\cite{Blas:2009qj,Blas:2010hb} gravity. Scalar and vector degrees of freedom in those theories induce dipole radiation that depends on the sensitivities of compact objects. So far, such sensitivities for strongly-gravitating objects have been calculated for NSs~\cite{Yagi:2013qpa,Yagi:2013ava}. One can repeat the analysis in~\cite{Yagi:2013qpa,Yagi:2013ava} by constructing slowly-moving BH solutions with respect to the vector field in these theories to extract BH sensitivities within the parameter space allowed from theoretical and observational constraints including GW170817~\cite{Gumrukcuoglu:2017ijh,Oost:2018tcv}. One can combine this with the orbital decay rate in these theories derived in~\cite{Yagi:2013ava} to estimate projected bounds on such theories from future BH-PSR observations. Another avenue for future work includes improving the analysis for bounding dCS gravity from BH-PSR observations. In this paper, we used the stellar mass BH quadrupole moment obtained within the slow-rotation approximation~\citep{DCSBoundSpin2,DCSBoundspin4} for BH-PSR systems with the BH dimensionless spin of unity, and thus the bounds should only be understood as order of magnitude estimates. One could improve this analysis by deriving the BH quadrupole moment valid for arbitrary spin. Such a goal may be achieved by using BH solutions with arbitrary spin in dCS gravity recently constructed numerically~\cite{Delsate:2018ome}. It would also be interesting to derive bounds on dCS gravity from the measurement of the advance rate of periastron for BH-PSRs. This is because Ref.~\cite{DCSKent} identified such an observable to be a useful post-Keplerian parameter when bounding dCS gravity from binary PSR observations. Typically masses are measured with e.g.~the advance rate of periastron and Shapiro time delay, and thus if GR tests are done with advance rate of periastron, one must use a different PPK parameter to derive the masses. This approach introduces more uncertainty to the bound because the mass measurement precision is reduced. Ideally, one needs to search for non-GR parameters and determine the masses simultaneously, instead of using the derived masses assuming GR is correct. | 18 | 8 | 1808.00080 |
1808 | 1808.07496_arXiv.txt | {We present a large-scale Bayesian inference framework to constrain cosmological parameters using galaxy redshift surveys, via an application of the Alcock-Paczy\'nski (AP) test. Our physical model of the non-linearly evolved density field, as probed by galaxy surveys, employs Lagrangian perturbation theory (LPT) to connect Gaussian initial conditions to the final density field, followed by a coordinate transformation to obtain the redshift space representation for comparison with data. We implement a Hamiltonian Monte Carlo sampler to generate realizations of three-dimensional (3D) primordial and present-day matter fluctuations from a non-Gaussian LPT-Poissonian density posterior given a set of observations. This hierarchical approach encodes a novel AP test, extracting several orders of magnitude more information from the cosmic expansion compared to classical approaches, to infer cosmological parameters and jointly reconstruct the underlying 3D dark matter density field. The novelty of this AP test lies in constraining the comoving-redshift transformation to infer the appropriate cosmology which yields isotropic correlations of the galaxy density field, with the underlying assumption relying purely on the geometrical symmetries of the cosmological principle. Such an AP test does not rely explicitly on modelling the full statistics of the field. We verify in depth via simulations that this renders our test robust to model misspecification. This leads to another crucial advantage, namely that the cosmological parameters exhibit extremely weak dependence on the currently unresolved phenomenon of galaxy bias, thereby circumventing a potentially key limitation. This is consequently among the first methods to extract a large fraction of information from statistics other than that of direct density contrast correlations, without being sensitive to the amplitude of density fluctuations. We perform several statistical efficiency and consistency tests on a mock galaxy catalogue, using the SDSS-III survey as template, taking into account the survey geometry and selection effects, to validate the Bayesian inference machinery implemented.} | \label{intro} The past few decades have witnessed the advent of an array of galaxy redshift surveys, with the state-of-the-art catalogues mapping millions of galaxies with precision positioning and accurate redshifts. The Sloan Digital Sky Survey (SDSS) \citep{sdss2000technical, sdss2009dr7, sdss2014dr10, sdss2015dr12} and the Six Degree Field Galaxy Redshift Survey (6dFGRS) \citep{6df2009dr3} are two notable examples. Future cutting-edge surveys from the Euclid \citep{euclid2011report, euclid2016missiondesign, euclid2016cosmology} and Large Synoptic Survey Telescope (LSST) \citep{lsst2008summary} missions, currently under construction, further highlight the wealth of galaxy redshift data sets which would be available within a five to ten year time frame. Sophisticated and optimal data analysis techniques, in particular large-scale structure analysis methods, are in increasing demand to cope with the present and upcoming avalanches of cosmological and astrophysical data, and therefore optimize the scientific returns of the missions. With the metamorphosis of cosmology into a precision (and data-driven) science, the three-dimensional (3D) large-scale structures have emerged as an essential probe of the dynamics of structure formation and evolution to further our understanding of the Universe. The two-point statistics of the 3D matter distribution have developed into key tools to investigate various cosmological models and test different inflationary scenarios. Various techniques to measure the power spectrum and several reconstruction methods attempting to recover the underlying density field from galaxy observations are described in literature \citep[e.g.][]{bertschinger1989recovering, bertschinger1991mapping, hoffman1994wiener, lahav1994wiener, fisher1995wiener, sheth1995constrained, webster1997wiener, bistolas1998nonlinear, schmoldt1999density, saunders2000interpolation, zaroubi1999wiener, zaroubi2002unbiased, erdogdu20042df, erdogdu2006reconstructed}, with the recent focus being on large-scale Bayesian inference methods \citep[e.g.][]{kitaura2008bayesian, kitaura2009cosmic, jasche2010fast, jasche2010bayesian, jasche2012bayesian, jasche2013methods, jasche2015matrix, jasche2018physical}. A formal and rigorous Bayesian framework provides the ideal setting to solve the ill-posed problem of inferring signals from noisy observations, while quantifying the corresponding statistical uncertainties. The potential of such Bayesian algorithms to jointly infer cosmological constraints, nevertheless, has not yet been exploited. We present, for the first time, a non-linear Bayesian inference framework for cosmological parameter inference from galaxy redshift surveys via an implementation of the Alcock-Paczy\'nski \citep[AP, ][]{alcock1979evolution} test. We extend the hierarchical Bayesian inference machinery of \textsc{borg} (Bayesian Origin Reconstruction from Galaxies) \citep{jasche2013bayesian}, originally developed for the non-linear reconstruction of large-scale structures, to constrain cosmological parameters. \textsc{borg} encodes a physical model for gravitational structure formation, yielding a highly non-trivial Bayesian inverse problem. This consequently allows us to reformulate the standard problem of present 3D density field reconstruction as an inference problem for initial conditions at an earlier epoch from current galaxy observations. \textsc{borg} builds upon the implementation of the Hamiltonian Monte Carlo (HMC) method \citep{neal1993probabilistic}, initially introduced in the \textsc{hades} (HAmiltonian Density Estimation and Sampling) algorithm \citep{jasche2010fast}, for efficiently sampling the high dimensional and non-linear parameter space of possible initial conditions at an earlier epoch. In this work, the conceptual framework is to constrain the comoving-redshift coordinate transformation and therefore infer the appropriate cosmology which would result in isotropic correlations of the galaxy density field. The key aspect of this application of the AP test consequently lies in its robustness to a misspecified model and the approximations therein, yielding a near-optimal exploitation of the model predictions, without relying on its accuracy in modelling the scale dependence of the correlations of the density field. Here, we employ Lagrangian Perturbation Theory (LPT) as a physical description for the non-linear dynamics and perform a joint inference of initial conditions, and consequently the corresponding non-linearly evolved density fields and associated velocity fields, and cosmological parameters, from incomplete observations. This augmented framework with cosmological applications is designated as \textsc{altair} (ALcock-Paczy\'nski consTrAIned Reconstruction). The paper is organized as follows. In Section \ref{ap_test}, the underlying principles of the AP test are outlined, followed by a description of the forward modelling approach and data model implemented in Section \ref{forward_modelling}. We then test the algorithm in Section \ref{results} on an artificially generated galaxy survey, with the mock generation procedure described in the preceding Section \ref{mock_generation}, by investigating its performance via statistical efficiency and consistency tests. In Section \ref{conclusion}, we summarize the main aspects of our work and discuss further possible extensions to our algorithm in order to fully exploit its potential in deriving cosmological constraints. In Appendix \ref{lpt_poisson_posterior}, we describe the LPT-Poissonian posterior implemented in this work, followed by the computation of the Jacobian of the comoving-redshift transformation in Appendix \ref{jacobian}. We provide a brief overview of the Hamiltonian sampling approach in Appendix \ref{hmc}, and follow up by deriving the required equations of motion in Appendix \ref{equations_of_motion}, with the numerical implementation outlined in Appendix \ref{numerical_implementation}. We subsequently describe how we increase the efficiency of our cosmological parameter sampler via a rotation of the parameter space in Appendix \ref{rotation_cosmo}. Finally, we outline the derivation of the adjoint gradient for a generic 3D interpolation scheme in Appendix \ref{adjoint_interpolation}. | \label{conclusion} We presented the implementation of a robust AP test that performs a detailed fit of the cosmological expansion via a non-linear and hierarchical Bayesian LSS inference framework. This forward modelling approach employs LPT as a physical description for the non-linear dynamics and sequentially encodes the cosmic expansion effect for joint inference of the cosmological parameters and underlying 3D density fields, while also fitting the mean density of tracers and bias parameters. In essence, this inference machinery explores the various cosmological expansion histories and selects the cosmology-dependent evolution pathways which yield isotropic correlations of the galaxy density field, thereby constraining cosmology. We demonstrated the application of our algorithm \textsc{altair} on an artificially generated galaxy catalogue, consisting of four subcatalogues, that emulates the highly structured survey geometry and selection effects of SDSS-III. We performed a series of statistical efficiency and consistency tests to validate the methodology adopted and showcased its potential to yield tight constraints on cosmological parameters from current and future galaxy redshift surveys. The main strength of our implementation of the AP test lies in its robustness to a misspecified model and its inherent approximations, thereby near-optimally exploiting the model predictions, without relying on its accuracy in modelling the scale dependence of the correlations of the field. Moreover, another key aspect of our approach, resulting from the robustness to a misspecified model, is that the cosmological constraints show extremely weak dependence on galaxy bias. This yields two crucial advantages. First, this is especially interesting as the lack of a sufficient physical description of this bias remains a potential limiting factor for standard approaches of cosmological parameter inference from such redshift surveys. Furthermore, this lack of sensitivity to the bias also implies that our method does not depend on the absolute density fluctuation amplitudes. This is therefore among the first methods to extract a large amount of information from statistics other than that of direct density contrast correlations, without relying on the power spectrum or bispectrum, thereby providing complementary information to state-of-the-art techniques. There is scope for further development of the \textsc{altair} framework, such as incorporating power spectrum inference, which is a highly non-trivial undertaking. We also intend to augment the current formalism to include the treatment of redshift space distortions, which is key for unbiased constraints on the cosmological parameters, and apply \textsc{altair} on state-of-the-art galaxy redshift catalogues for cosmological inference. | 18 | 8 | 1808.07496 |
1808 | 1808.08391_arXiv.txt | Since the CoRoT and Kepler missions, the availability of high quality seismic spectra for red giants has made them the standard clocks and rulers for Galactic Archeology. With the expected excellent data from the TESS and PLATO missions, red giants will again play a key role in Galactic studies and stellar physics, thanks to the precise masses and radii determined by asteroseismology. The determination of these quantities is often based on so-called scaling laws, which have been used extensively for main-sequence stars. We show how the SOLA inversion technique can provide robust determinations of the mean density of red giants within 1 per cent of the real value, using only radial oscillations. Combined with radii determinations from Gaia of around 2 per cent precision, this approach provides robust, less model-dependent masses with an error lower than 10 per cent. It will improve age determinations, helping to accurately dissect the Galactic structure and history. We present results on artificial data of standard models, models including an extended atmosphere from averaged 3D simulations and non-adiabatic frequency calculations to test surface effects, and on eclipsing binaries. We show that the inversions provide very robust mean density estimates, using at best seismic information. However, we also show that a distinction between red-giant branch and red-clump stars is required to determine a reliable estimate of the mean density. The stability of the inversion enables an implementation in automated pipelines, making it suitable for large samples of stars. | \label{sec:intro} Red giants play a key role in stellar physics. Since the detection of mixed modes in their oscillation spectra \citep{DeRidder} thanks to the CoRoT \citep{Baglin} and \textit{Kepler} \citep{Borucki} missions, they are at the origin of multiple questions on the reliability of our depiction of stellar structure and evolution \citep{Mosser,Deheuvels}. The availability of thousands of high quality seismic spectra for these stars led to their use as the standard clocks and rulers \citep{MiglioClocks} for Galactic Archeology \citep{MiglioPopulations,Anders2017RG}. Today, they are used as tracers of the structure and chemical evolution of our Galaxy \citep{AndersRG}. New accurate data for these stars are also expected to be delivered by the TESS and PLATO \citep{Rauer} missions, which will play a key role in Galactic studies \citep{Miglio2017}. These successes originate from the ability of asteroseismology to provide precise masses and radii for a large number of stars. The seismic determination of these quantities is often based on so-called scaling laws, which have also been used extensively for main-sequence stars. However, while the precision of these determinations is excellent, due to the high precision of the space-based photometry data, their accuracy is far from perfect. Multiple studies \citep{Gaulme2016,Brogaard2016,Rodrigues2017, Viani2017, Brogaard2018} have shown that they could lead to inaccurate results. From a physical point of view, their limited accuracy is not surprising as they do not fully exploit the information contained in the seismic spectra. Therefore, providing a more robust way of determining the mean density of the observed targets using seismology is required so that, using constraints from the second Gaia data release, more accurate masses can be determined for thousands of stars. These accurate masses will help with dissecting the structure of the Galaxy, thus providing new insights on its evolution and formation history. In addition to this potential, the determination of accurate fundamental parameters of red giants in stellar clusters is also crucial to constrain the mass loss rate on the red giant branch, a still uncertain key phenomenon of stellar evolution \citep{MiglioClusters, HandbergRG}. In this study, we will show how the adaptation of the SOLA inversion technique for the mean density developed by \citet{Reese} used on the radial oscillations of red giant stars can provide more robust determinations of the mean density than values obtained from the fitting of the average large frequency separation or the usual scaling laws. In section \ref{sec:inversion}, we briefly recall the principles of the inversion techniques. In section \ref{sec:numex}, we test the inversion in various numerical exercises, using artificial targets on the red giant branch (hereafter denoted RGB), in the red clump, and an RGB target including an averaged $3$D atmosphere model for which the frequencies are computed using adiabatic and non-adiabatic oscillation codes to test various surface effects correction. In section \ref{sec:ObsBin}, we apply our method to a subsample of eclipsing binaries studied by \citet{Gaulme2016} and \citet{Brogaard2018}. This is then followed by a conclusion. | \label{sec:Conclu} In this study, we have demonstrated the feasibility and robustness of seismic inversions of the mean density of red giant and red clump stars using only a few observed radial modes, providing an extension to the framework of the approach initially applied to main-sequence solar-like stars \citep{Reese, Buldgen}. We have started by introducing the approach to the inverse problem in Sect. \ref{sec:inversion} and applied it in extensive numerical tests using calibration techniques based on effective temperature and luminosity in Sect. \ref{sec:TeffL} as well as seismic constraints \ref{sec:SismoCons}. We have analysed the possibility of carrying out inversions for the mean density both for red giant branch and clump stars, as well as trying to determine the mean density of a misidentified clump star. In the last case, the inversion proved to be inaccurate and thus the method provided here is only valid for unambiguously identified red giant branch and clump stars. This last point also proves the need for a reliable reference model before carrying out the inversion. This weakness has also been observed for other numerical exercises were the so-called cross-term errors could contribute very significantly to the total error budget of the inversion and even dominate other errors. This is radically different from inversions on the main-sequence and is due to both the small number of modes and the large radial extent of the ionization zones, leading to more widespread differences in $\Gamma_{1}$ from one model to the other. In addition to testing the use of seismic constraints, we also carried out in Sect. \ref{sec:SurfEff} inversions for an artificial target from \citet{Sonoi} including an atmospheric model from an averaged hydrodynamical simulation. For this target, we used both adiabatic and non-adiabatic frequencies to carry out the inversion. These tests illustrated the impact of surface effects and the importance of correcting them. It should be noted that none of the current methods seemed to work perfectly, as error compensations were sometimes observed and that directly implementing the corrections as free parameters in the SOLA method provided inaccurate fits of the target function of the inversion. In Sect. \ref{sec:ObsBin}, we carried out mean density inversions on observed red giants in eclipsing binary systems from the studies of \citet{Gaulme2016} and \citet{Brogaard2018}. We first showed that using individual frequencies, corrected for surface effects using the approach of \citet{Ball1}, lead to a much better agreement in terms of mass and radius than simply using scaling relations or global seismic indices. Furthermore, we showed that the mean density inversions could provide either further small corrections to the mean density obtained from forward modelling or an additional verification step. Combining the inverted mean density to the radii determinations from eclipses provided an overall good agreement in terms of masses for these stars. However, these test cases also demonstrate the importance of classical constraints for the modelling of red giants, since the most significant improvement came from using the dynamical radii values. Indeed, the mean density of these stars was already very accurately determined through forward modelling. A couple of conclusions can be drawn from this last point. Firstly, that pure seismic modelling using only radial modes might lead to degeneracies since the frequencies will be mostly sensitive to the mean density of the star. Secondly, directly using the individual frequencies might lead to underestimated uncertainties and the seismic constraints might dominate the classical constraints. Therefore, using directly an inverted value for the mean density, alongside a precise and accurate value for the luminosity, the $\left[\mathrm{Fe/H}\right]$ and the radius might provide a more direct and balanced approach to the modelling of red giant stars. Other seismic constraints could also be used, such as the asymptotic period spacing or ratios of radial oscillation frequencies as is done for Cepheids and other classical pulsators. The application of such approaches to an extended set of eclipsing binaries will provide a unique opportunity to test the reliability of seismic modelling and the importance of classical constraints. In the near future, using such approaches alongside constraints from the second Gaia data release will help better understand the properties of red giants. Providing more accurate masses is indeed crucial to determine the properties of various stellar populations in the Galaxy but also for example to pinpoint the properties of additional mixing at the base of the convective envelope, manifesting itself through the so-called RGB bump \citep{Alongi91,Cassisi2011,Khan2018}. In stellar clusters, accurate masses could also be used to characterize mass loss on the red giant branch \citep{HandbergRG}, one of the major issues in current stellar evolution models. | 18 | 8 | 1808.08391 |
1808 | 1808.01086_arXiv.txt | We solve the Combustion adiabat (CA) or the Chapman-Jouget adiabat equation to study the phase transition (PT) of a neutron star (NS) to a quark star (QS). The hadronic matter and quark matter equation of states are used to calculate the matter velocities on either side of the shock front. The CA with the hadronic matter as an input is solved to obtain the corresponding quark matter values. The maximum of the quark pressure is reflected in the retracing of the path in the CA curve. The downstream quark pressure maximum implies towards a maximum mass limit of a phase transformed QS which is different from the regular mass limit of an ordinary QS. Further, the characterization of velocities suggest that the PT from NS to QS is not always feasible from the center of the star. The possible mode of combustion in NSs is likely to be a slow deflagration in most of the low and intermediate density range. The result is crucial and emphasizes on the fact that PT in NSs does not always starts from the center and sometimes a NS does not suffers a PT at all. | Relativistic shocks are very common in astrophysical stellar objects and are associated with stellar winds, supernovae remnants, radio jets, accretion onto compact objects and phase transition (PT) in neutron stars (NSs). The results of the relativistic Rankine-Hugoniot (RH) jump condition is well-studied \citep{landau} in connection with shocks, with their solution being still open for discussion. The RH condition can be used to derive a single equation relating the matter quantities across the discontinuity, known as Taub adiabat (TA) or shock adiabat relation \citep{taub}. Thorne \citep{thorne} showed that the TA method of solution could be carried over from non-relativistic to relativistic shocks. In the TA, the initial and final states are the same functions of pressure, energy density and density (same EoS), and hence they lie on the same curve. However, the form of the equation still holds if the initial and final states are not the same function of pressure, energy density, and density (different EoS). This happens because of the difference in the chemical energy of the initial and final state. Therefore, there is combustion from the initial state to the final state. Due to the difference in the EoSs the final state curve shifts from the initial state curve. The relation connecting the initial and final state is called the combustion adiabat (CA) or the Chapman-Jouget adiabat. Several authors have used the relativistic RH condition to study various problems in astrophysics, but the study of CA in an astrophysical scenario is insufficient. Relativistic shock phenomena have been widely discussed in recent years with regards to their connection to the PT in NSs. NSs are thought to be the best candidates depicting a phase of strongly interacting deconfined matter. They serve as natural laboratories to study the low temperature and high density (or high baryon chemical potential) regime of the QCD phase diagram. Since the proposition that strange quark matter can be the stable configuration at such high densities \citep{itoh,bodmer, witten}, there has been considerable effort to investigate this possibility in astrophysical literature. One of the lines of investigation assumes it to be a simple first order PT from confined hadronic matter (HM) to deconfined quark matter (QM). Accretion process in NS can trigger this PT \citep{cheng2} or nucleation via seeding can also occur \citep{alcock}. The process of phase transition of matter at extreme astrophysical densities is very complicated and highly debated. Olinto \citep{olinto} was one of the earlier physicists to study the combustion process. She viewed it as slow combustion process where an excess of down quark gets converted to strange quark via the weak interaction process. Collins and Perry \citep{collins} argued that it is rather a two-step process, where the initial nuclear matter gets converted to final strange quark matter (3-flavour (3f)) via an intermediate step of 2-flavour (2f) quark matter. Horvath and Benvenuto \citep{horvath} argued that the convective instability could turn slow combustion (deflagration) process to a detonation. Cho et al. \citep{cho} used the hydrodynamic RH jump conditions to conclude that weak detonation is the correct mode of combustion. However, Tokareva et al. \citep{tokareva} suggested that the detonation modes can also be possible whereas Lugones et al. \citep{lugones} strongly advocated that the actual mode of combustion is strong detonation. Bhattacharyya et al. \citep{mallick1} introduced two-step combustion involving the hydrodynamic jump equations to study the first step and slow weak combustion for the second step. However, Drago et al. \citep{drago} argued the combustion mode to be a slow deflagration involving a mixed phase region. Recently Prasad \& Mallick \citep{mallick-apj} did a dynamical evolution of the phase transition front and found that the PT takes not more than $50 \mu$s to occur. Niebergal et al. \citep{neibergal} numerically solved for the front velocity involving the hydrodynamic equations along with neutrino emission and weak interaction of down to strange quark conversion and found that the combustion velocity of the weak process is very high. Recently Furusawa et al. \citep{furusawa1,furusawa2} did a detailed discussion on scenarios involving shock induced and diffusion induced PT. In our analysis, we are dealing with shock-induced PT where the shock initiates a PT (or combustion) in NS. Instead of using the RH jump condition we employ the CA equation to study the PT of HM to QM in neutron stars. The CA is easily solvable as it a single equation devoid of any velocity terms. The paper is arranged as follows: In section 2 we give the details of the EoS that we have used in our work. Section 3 deals with the study of the CA where we employ HM and QM EoS in the CA to show our results. Finally, in section 4 we summarize our findings and conclude from them. | To summarize, we have employed CA as a tool for studying the PT from NS to QS. Studies using the RH equations have been done earlier, where three equations are solved for three unknowns. The velocities occur as an additional parameter for these equations. However, this CA technique is independent of velocity, and only the upstream and downstream EoSs are enough to study their PT. Our main results are shown using two HM EoS (NL3 and PLZ parameter setting) and a QM EoS (with bag pressure ($140$ MeV)$^4$). For comparison, we have also shown results with a single HM EoS (PLZ) and for three QM EoSs having three different bag pressures (Q-120, Q-140, and Q-160). EoS having polytropic form has also been used to check our results. Using HM EoS for the upstream curve and the QM EoS for the downstream burnt curve we solve the CA to obtain the respective downstream quantities. The whole of the downstream CA curve can be generated by varying the initial input upstream parameters. Solving the CA leads to a maximum pressure in the downstream QM state. After the maximum, the CA retraces its path. The retracing nature of the burnt QM CA curve is quite robust and is obtained for both hard and soft EoSs. Although the upstream HM pressure always increases, the burnt QM pressure has a maximum. The retracing nature of the CA arises due to the occurrence of the maximum pressure of the burnt QM. This nature is not only limited to EoS of HM or QM but can also be extended to polytropic EoS which is sometimes used to calculate NS properties. The above analysis of the CA indicates at a maximum limit on the QS (HS/SS). This mass limit is smaller than the Chandrasekhar mass limit for QSs. The study of the velocity to characterize the PT shows that the PT of NS to QS does not always start at the center of the star. For some particular central densities, the PT cannot trigger from the core, and the combustion begins at some radial point in the star located at some distance from the center. Observing the criterion for detonation and deflagration from the velocity of the respective phases, we found that the velocity of the NM in the low and intermediate density range is higher than that of the QM, which is the condition for a deflagration. For very high densities there is a possibility of detonation. Also, there is a considerable density range for which velocity is either zero or becomes unphysical. We can hence conclude that for these particular densities at the star center the PT process cannot commence from the star center. We can also conclude that the burning process at the star center most likely starts as a deflagration process. For more massive stars the detonation process is a possible mode of combustion, however, from the pressure curve we can conclude that the massive star undergoing PT via detonation gives rise to a comparatively small QS. The results from the pressure maximum, the physical velocities, we get a density range where the PT is more likely. The stars outside this range are either not very likely to suffer PT or the PT in such stars cannot commence from the center of the star. We must emphasize here that in this study we are not performing any dynamical calculations nor we are doing any micro-physical analysis. Our results are purely based on the hydrodynamic study of the initial and final EoSs. Although, the retracing nature is a global phenomenon, the quantitative value of the results dependent on the choice of EoS sets we choose. The details of the micro-physics lie in the EoS, and we do not study them here. Also, the quantitative region of detonation, deflagration and the value of maximum mass limit inferred by the restricted density range for PT depends on the choice of the EoS employed. The maximum of the downstream variables and the retracing of the CA curve is quite robust. It seems that this is a characteristic of the CA and it could be a global property and not only related to astrophysical scenarios. Such nature of the CA can also be checked for ideal and realistic EoSs which can be verified in high energy experiments. Also the present study gave rise to some interesting phenomena and results which should be supplemented with dynamic studies. Our present effort is towards this regard. | 18 | 8 | 1808.01086 |
1808 | 1808.01615_arXiv.txt | {We discuss some implications of the recently suggested Swampland conjecture $\frac{|\nabla V|}{V}\gtrsim c\sim 1$, together with a previous one $\Delta \phi \lesssim 1$. We list some implications for particle phenomenology and the Early Universe. The most intriguing implication of the conjecture could be a significant shift in allowed inflationary models, if not ruling out slow-roll (single field) inflation altogether. The tension of inflation and the conjecture does not only regard the amplitude of the tensor spectrum, but also its tilt, as $c\gtrsim 1$ implies both a yet unobserved tensor to scalar ratio, and an enhancement of the observed scalar power spectrum on large scales in discord with current data that favors a suppression on these scales. Scalar fields are abundant in theories of quantum gravity. Considering a second scalar field, its dynamics are dictated by the relation between its mass, $m$, and the Hubble parameter, $H$, at different epochs in the history of the Universe. This scalar field, a drainon, fulfills the conjecture draining up the swampland. For Inflation, this drainon requires a modest hierarchy compared to the inflaton. For the rest of the thermal history of the Universe, the drainon can be a coherently oscillating scalar field strengthening the case of Dark Matter candidates of that sort. % }% | Based on lingering debates about the consistency of de Sitter (dS) space and Quantum Gravity, it has been suggested that for any consistent theory of Quantum Gravity there exists an inequality \cite{Obied:2018sgi}: \be \textbf{C1:} \quad \frac{M_{pl}|\nabla V|}{V}\gtrsim c\sim 1 , \quad V>0. \ee The authors in \cite{Obied:2018sgi} give some supporting evidence for this claim, and in return it has been used to place bounds on the validity of inflationary and quintessence models \cite{Agrawal:2018own}. In the literature there are known counter examples to the conjecture, \cite{Kachru:2003aw, Brustein:2004xn,Balasubramanian:2005zx,Westphal:2006tn,BenDayan:2008dv,Rummel:2011cd}. However, there always seem to be lingering doubts about the validity of these constructions, for example, \cite{Danielsson:2018ztv,Sethi:2017phn}. Therefore, it makes sense to consider the implications of the conjecture on phenomenology, and whether the conjecture is true or may have to be revised. Let us stress that the argument is not limited to the specific point, but rather some domain in field space. Therefore, the argument states that within a domain, $\Delta \phi$, the inequality holds. Furthermore, earlier conjectures based on the weak gravity conjecture and some explicit examples suggest \be \textbf{C2:}\quad \Delta \phi\lesssim M_{pl}, \ee in field space \cite{Grimm:2018ohb}. Beyond $\Delta \phi\sim M_{pl}$ additional light states are expected to appear in the spectrum of the theory invalidating the analysis. Taken each criterion separately, phenomenological implications seem rather minimal. It may cast doubts on the validity of some inflationary models, but the literature is filled with constructions that fulfill one of the criterion or the other. However, taking into account both criteria seems to severely constrain low energy effective field theories. Various suggestions and interpretations have been put forward \cite{Denef:2018etk,Ghosh:2018fbx}. The interpretations involve several logical paths. \begin{itemize} \item Revising C1, for example, taking $c$ to be of the order of the slow-roll parameters $c\sim O(0.01)$. \item Revising $\Lambda CDM$ and/or Inflation, pointing towards quintessence as the cause of present acceleration and bouncing models rather than inflation as the mechanism that fits the known CMB data. \item Disproving the conjectures via explicit constructions \item A combination of the above. \end{itemize} In this note, we analyze two approaches. First, we stick to \textit{strictly} single field analysis and discuss the implications of the criteria for Higgs Physics and Inflation. Regarding the Higgs field, the criteria are in contradiction with the standard picture of spontaneous symmetry breaking and the electroweak phase transition. Regarding Inflation, we show that even C1 by itself, for any field range leads to interesting bounds on the tensor to scalar ratio $r$. The bound on $r$ is due to the correspondence between the slow-roll parameters and the CMB observables. Second, the spectrum of any realistic fundamental theory is \textit{never} a single field. The Standard Model fields are an immediate example. Even if single field inflation was realized in nature, fundamental theories and in particular string theory predict the existence of additional scalar fields, but their energy density and dynamics are negligible compared to the inflaton, which is why we can neglect them in deriving predictions of single field model or in analyzing inflationary dynamics. Specifically, the curvaton does not affect the inflationary dynamics, but its dynamics may have generated the observed CMB scalar spectrum \cite{Bartolo:2002vf}. Hence, to fulfill both criteria we have to take into account all fields. We show how a second scalar field, a "drainon", allows to fulfill both criteria, hence draining the swampland. We apply the drainon idea for inflation and the thermal history of the Universe after inflation. For the evolution of the Universe after inflation, the criteria plus simplicity point to the drainon being a coherently oscillating scalar field, strengthening the case for such Dark Matter models. For inflation, the drainon is stuck at some point of the potential like a curvaton, but has no observable consequences. \textbf{Note added:} While the preprint was being finalized \cite{Denef:2018etk} appeared which also showed how the Higgs potential is in contradiction with the dS swampland condition C1. | The conjectures C1 and C2, and their predecessor the weak gravity conjecture, seem to pose an interesting challenge for Inflation, particle phenomenology and the Cosmological Constant. We have shown that the main challenge lies when assuming a model of a strictly single field, while considering a more realistic model allows to introduce a drainon that drains most of the swampland. Given an oscillating massive scalar, the problem of the present acceleration is pushed into the distant future, where $H<10^{-33} eV$, and the possibility of inflation is tenable given a drainon with a logarithmic potential. The underlying motivation of the recent C1 and C2 conjectures is the difficulty in constructing a dS metastable vacua or inflationary model from the basic ingredients of a given string theory compactification. As such, having to accommodate a drainon field could be even more challenging. However, if model builders do mange to construct metastable dS or inflation in string theory explicitly, then the conjectures become irrelevant anyway and the drainon unnecessary. Given the weak evidence of the conjecture C1, the knowledge of the actual value of $c$ and the fact that a drainon seems to work rather well even for inflation, it seems premature to deviate from Inflation and $\Lambda$CDM that until now have been such a successful match to the data. | 18 | 8 | 1808.01615 |
1808 | 1808.06203_arXiv.txt | The laws of geometric optics and their corrections are derived for scalar, electromagnetic, and gravitational waves propagating in generic curved spacetimes. Local peeling-type results are obtained, where different components of high-frequency fields are shown to scale with different powers of their frequencies. Additionally, finite-frequency corrections are identified for a number of conservation laws and observables. Among these observables are a field's energy and momentum densities, as well as several candidates for its corrected ``propagation directions.'' | Nearly all astronomical observations involve, fundamentally, measurements of electromagnetic or (more recently) gravitational radiation. However, these waves carry with them an imprint of the spacetime through which they travel. The spacetime geometry provides a kind of ``transfer function'' that relates the intrinsic properties of a source to its radiated fields. Such relations must be understood if an object's properties are to be accurately inferred from distant measurements of its fields. If a source has already been characterized, its radiation might instead be used to probe the intervening geometry, and thus the matter which contributes to it---matter which might not be bright enough to observe directly. For these reasons and others, gravitational lensing has become a standard tool with which to extract information from astronomical observations. Much of the theory of gravitational lensing which is used in practice may be viewed as an elaboration on the particle-like laws of geometric optics: Light travels along null geodesics, intensity variations are determined by the changing cross-sectional areas of ray bundles, and polarization vectors are parallel transported. These simple statements beget a remarkable variety of applications \cite{SchneiderEhlers,WambsganssRev, Perlick, BartelmannRev}. However, the laws of geometric optics are an approximation. Electromagnetic fields are more properly described as solutions to Maxwell's equations, and gravitational waves as solutions to Einstein's equation. While the full complexities of these equations may often be ignored, there are exceptions. For example, it is well-known in ordinary optics \cite{CausticsBook, BornWolf, ThorneBlandford} that the geometric approximation breaks down completely at caustics---a result which has also had astrophysical implications \cite{SchneiderEhlers, BlandfordCaustics, Turyshev2017a}. In other contexts, wave-optical corrections may be small but still detectable, in which case they might supply information which is different---and therefore complementary to---that which can be learnt from geometric optics alone. Wave-optical effects may be viewed as frequency-dependent corrections to the frequency-independent laws of geometrical optics. Apparent source locations, intensities, phases, and polarization states might all depend on the frequencies at which a source is observed. Any such quantity measured at a sufficiently-high characteristic frequency $\omega$ may be viewed as a geometric-optics result plus relative corrections which scale like, e.g., $\omega^{-1}$. Somewhat more precisely, these corrections scale like $(\omega \ell)^{-1}$, where $\ell$ is a relevant lengthscale. Several lengthscales may be present simultaneously and different ones can be relevant for different observables. In simple cases, $\ell$ might represent a notion of distance between a source and its observer: That geometric optics breaks down at caustics may be understood in this context by noting that the ``source-centered area distance'' $r_\mathrm{a}$ goes to zero at caustics and $\ell \sim r_\mathrm{a}$ for some contributions to some observables. More generally, $\ell$ can be a nontrivial composite of different lengthscales. For example, some corrections associated with fields of mass $\mu \neq 0$ which are radiated by a source at affine distance $r$ can involve the lengthscale $\ell \sim (\mu^2 r)^{-1}$; fractional corrections to geometric optics grow with distance for massive fields. More generally (and even for massless fields) a relevant $\ell$ might be a highly nontrivial nonlocal combination of different lengthscales---including those associated with the spacetime geometry and with details of the particular field under consideration. A systematic development of the underlying theory is thus required in order to understand precisely when such effects might be interesting. This paper begins on the path to such a development. More directly, the purpose here is to provide general expressions which allow the propagation of high-frequency scalar, electromagnetic, and gravitational waves to be characterized in general spacetimes. While the basic equations governing geometric optics and its corrections have been discussed elsewhere \cite{EhlersGeoOptics, Anile, Isaacson1} from a general spacetime perspective, very few of their higher-order consequences appear to have been explored. Some discussions which do go beyond geometric optics have appeared in various contexts, although most of these have employed a different ``pseudo-Euclidean'' approach which is restricted to weakly-curved spacetimes \cite{SchneiderEhlers, Nakamura1998, Nakamura1999, Takahashi2006, gravGrating}. The discussion here is intended to be largely self-contained, and therefore begins by reviewing the equations which govern geometric optics and its corrections. Mathematically, these equations transform the partial differential equations which control the underlying fields into a hierarchy of algebraic constraints and ordinary differential equations along null geodesics. These are used to derive wave-optical corrections to field strengths, curvature perturbations, stress-energy tensors, and conservation laws---in arbitrary spacetimes and for arbitrary polarization states. Several types of ``propagation direction'' are identified and discussed. For some such definitions, multiple directions can arise simultaneously; these experience relative corrections which scale like $\omega^{-1/2}$ instead of, e.g., $\omega^{-1}$, implying that they are particularly sensitive to wave-optical effects. Frequency dependencies of the different tensorial components of electromagnetic and gravitational waves are determined as well, resulting in what are essentially local peeling results. Throughout, we emphasize connections between the various types of fields considered here. When, for example, can aspects of an electromagnetic problem be reduced to those of an effective scalar problem? \noindent \textit{Notation}---Sign and index conventions follow those of Wald \cite{Wald}. Units are used in which $G=c=1$ and the number of spacetime dimensions is fixed at four. In several cases, a complex field is considered despite that it is only its real component which is considered to be physical. These fields are distinguishing by using an upper-case symbol to denote the real quantity and a lower-case one for its complex counterpart; $F_{ab} = \Re f_{ab}$, for example. | We have derived a number of general features of high-frequency scalar, electromagnetic, and gravitational waves propagating on curved background spacetimes, focusing on observables, physical intuition, and also relations between these different types of fields. However, no specific applications were considered. The purpose has been instead to set the stage for further exploration. While it would be straightforward to use the results presented here to compute corrections to geometric optics in various scenarios, subsequent papers in this series will take a more foundational approach. Two basic questions will be addressed before considering the details associated with any specific systems: First, how do changes in the background metric affect observables? General invariance properties of the underlying equations will be shown to provide powerful tools with which to address this question. Second, we ask how the measured properties of a radiated field can be related to intrinsic properties of its source. Alternatively, how should initial data be specified for the various transport equations? Although the space of possibilities is large in general, gravitational lensing is typically concerned with compact sources. In this context, the initial data problem simplifies considerably. We shall discuss how this occurs and how the relevant data can be related to a source's intrinsic properties. \appendix | 18 | 8 | 1808.06203 |
1808 | 1808.01109_arXiv.txt | We report the discovery of long-period radial velocity (RV) variations in six intermediate mass K giant stars using precise RV measurements. These discoveries are part of the Search for Exoplanets around Northern Circumpolar Stars (SENS) survey being conducted at the Bohyunsan Optical Astronomy Observatory (BOAO). The nature of the RV variations was investigated by looking for photometric and line shape variations. We can find no variability with the RV period in these quantities and conclude that RV variations are most likely due to unseen sub-stellar companions. Orbital solutions for the six stars yield orbital periods in the range 418 -- 1065 d and minimum masses in the range 1.9 -- 8.5 $M_{J}$. These properties are typical on planets around intermediate mass stars. Our SENS survey so far has about an 8\% confirmed planet occurrence rate, and it will provide better statistics on planets around giant stars when the survey is completed. | \label{sec:intro} The measurement of stellar radial velocity (RV) is one of the most effective techniques employed in the search for exoplanets. At the Bohyansan Optical Astronomical Observatory (BOAO), we have been conducting an exoplanet search program around late-type giant stars since 2004. This program has made contributions to both exoplanet and asteroseismic studies around K giant stars \citep{2008JKAS...41...59H,2010A&A...509A..24H,2011A&A...529A.134L,2012A&A...546A...5L,2012A&A...548A.118L,2014A&A...566A..67L} and exoplanet detection around G giant stars \citep{2009PASJ...61..825O, 2012PASJ...64...34O, 2013PASJ...65...85S}. In 2010, we began a new program, the Search for Exoplanet around Northern circumpolar Stars (SENS; \citealt{2015A&A...584A..79L}). The main goal of SENS is to observe stars that are accessible year-round in order to have better sampling for our targets and thus increase the planet detection efficiency. The stars of SENS were selected from the \textit{HIPPARCOS} catalog with visual magnitudes of 5.0 $<$ $m_{v}$ $<$ 7.0 and color indices of 0.6 $<$ $\textit{B -- V}$ $<$ 1.6. The original SENS sample consist of 224 stars -- 5\% dwarfs, 40\% giant stars, and 55\% unclassified stars. From SENS survey, we detected periodic RV variations around roughly twenty G, K, and M giant stars. Among them, \mbox{HD 104985} \citep{2003ApJ...597L.157S}, 11 Ursae Minoris \citep{2009A&A...505.1311D}, \mbox{HD208527} \& \mbox{HD 220074} \citep{2013A&A...549A...2L}, \mbox{HD 11755}, \mbox{HD 12648}, \mbox{HD 24064} and 8 Ursae Minoris \citep{2015A&A...584A..79L}, \mbox{HD 36384}, \mbox{HD 52030}, and \mbox{HD 208742} \citep{2017ApJ...844...36L} were later shown to host planetary companions. In this paper, the observational strategy and data reduction are summarized in Section 2. Section 3 describes the stellar properties and analysis of each host stars. In Section 4, orbital solutions are described in detail. Some possible causes of the RV variations are examined in Section 5. The discussion about the results is presented in Section 6. | \label{sec:dis} We have found long-period RV variations of six K-giant stars. The stars show no variations in the line shapes as measured by the spectral line bisectors. HD~44385 does show a weak signal in the bisector curvature measurements (BVS), but this does not seem to be significant (FAP $\ge$ 0.1). We note that a lack of bisector measurements is not proof of the planetary nature. High quality bisector measurements are difficult to make as these require high spectral resolution and high S/N data. In general these are of much lower quality than the RV measurements. If we found bisector variations with the RV period, then the planet hypothesis is refuted. On the other hand, a lack of bisector variations is not sufficient proof of the existence of the planet. It could well be that a phenomenon produces measurable RV variations, but small bisector variations that are difficult to measure. The stars also seem to show a lack of variations in the \textit{HIPPARCOS} photometry. Only one star, HD~106574, shows a weak peak in the L-S periodogram at the RV period. Again, this peak seems to be of low significance. However, the lack of photometric variations is only suggestive as the \textit{HIPPARCOS} measurements were not contemporaneous to our data. HD~44385, HD~106574, and HD~118904 have larger RV scatters than those predicted by the \citet{1995A&A...293...87K} relationship. Later spectral type stars, such as HD~106574 and HD~118904, show larger RV scatters than those given in \citet{2005PASJ...57...97S,2006A&A...454..943H}. \citet{2012A&A...548A.118L} also have detected an exoplanet with similar orbital parameters and rms of the RV residuals. Some more exoplanets discovered using BOES have shown large RV scatters \citep{2013A&A...549A...2L,2014A&A...566A..67L,2014JKAS...47...69L,2015A&A...584A..79L,2015A&A...580A..31H,2018AJ....155..120H}. RV scatter in giant stars may have its origin not only from the stellar pulsations but also from the stellar activities. \citet{1993ApJ...413..339H} found RV variations in $\alpha$ Boo with a period of 233 d and an amplitude of $\sim$200 m s$^{-1}$. This RV period was the same period found in the He I 10830 variations in this star by \citet{1987ApJS...65..255L}. He I 10830 is a chromospheric activity indicator. In this case, large RV variations are clearly due to the stellar activity. Activity in giant stars is poorly understood and it may be that these are often not accompanied by variations in the ``classic'' indicators of stellar activity. We do not know the timescales or RV amplitudes of such activity jitter for these stars, which may have contributed to the RV scatters seen in some of our stars. The exact cause of large RV scatters seen in our stars is yet to be understood by further study. Rotating stars with surface features can also exhibit periodic RV variations which can be mimic a ``planetary like'' signal. Our stars appear to be relatively inactive as shown by an absence of emission in the core of Ca II H lines. Another check on rotational modulation can come from estimates of the rotational period of the star. From $v_{\rm{rot}}$ sin $i$ and $\textit{$R_{\star}$}$ (Table~\ref{tab:ste}) we estimate upper limits on the rotational period of 354.1 $\pm$ 72.0 days for \mbox{HD 44385}, 444.7 $\pm$ 122.0 days for \mbox{HD 97619}, 508.9 $\pm$ 152.1 days for \mbox{HD 106574}, 624.0 $\pm$ 265.7 days for \mbox{HD 118904}, 629.6 $\pm$ 125.7 days for \mbox{HD 164428}, and 267.4 $\pm$ 64.1 days for \mbox{HD 202432}. \mbox{HD 118904} and \mbox{HD 164428} have estimated rotation periods comparable to the RV periods. In the case of \mbox{HD 44385}, \mbox{HD 97619}, \mbox{HD 106574} and \mbox{HD 202432} are the maximum rotational periods significantly larger than the RV periods. For these four stars our estimated rotational periods provide further support that we are not seeing rotational modulation. Another explanation for the RV variations in the six stars is a new, unknown form of stellar oscillations. One possibility is oscillatory convective modes. These have been proposed to explain the long-period variables \citep{2015MNRAS.452.3863S}, stars that are more evolved than the K giant stars of our work. Coincidentally, only one of stars, HD 164428, is rather evolved with stellar radius larger than 20 $R_\odot$. However, given that we know so little about long-period oscillations in K giant stars it seems that at the present time the most likely explanation for the RV variations in our stars is Keplerian motion by planetary companions. All of our targets are evolved intermediate mass stars. If the variations are due to planetary companions our detections add to the sample of exoplanets around stars more massive than the sun. Two of our stars have masses $\ge$ 1.5 $M_\odot$. Of the approximately 2200 exoplanets with good orbits and planet mass determinations, less than 5\% of the host stars have masses greater than 1.5 $M_\odot$ (source: exoplanets.org). Exoplanets around host stars with M $\ge$ 2 $M_\odot$ have a median $m$ sin $i$ of 2.7 $M_{J}$ and a median orbital period of $\sim$ 400 d. Thus the companions to our six K giants have properties that are typical for planets around massive stars: massive giant planets (1.9 -- 8.5 $M_{J}$) with orbital periods of several hundreds of days. Although the SENS survey is not yet complete, at this stage we can still make a rough estimate of the planet frequency of our sample. We have found periodic variations in 31 of our sample of 224 stars. Among these, 17 stars have confirmed planets which is a planet occurrence rate of about 8\%. If all the RV variations are planetary in nature, then the occurrence rate can be as high as $\approx$ 15\%. This is largely in line with the expectation that $\sim$10\% giant stars have planetary companions \citep{2015ApJ...798..112M}. A more detailed analysis of the statistics can be made once the SENS survey is completed. | 18 | 8 | 1808.01109 |
1808 | 1808.07425_arXiv.txt | The Mexican Space Weather Service (SCiESMEX in Spanish) and National Space Weather Laboratory (LANCE in Spanish) were organized in 2014 and in 2016 respectively to provide space weather monitoring and alerts, as well as scientific research in Mexico. In this work, we present the results of the first joint observations of two events (22 June, 2015, and 29 September, 2015) with our local network of instruments and their related products. This network includes the MEXART radio telescope (solar flare and radio burst), the Compact Astronomical Low-frequency, Low-cost Instrument for Spectroscopy in Transportable Observatories (CALLISTO) at MEXART station (solar radio burst), the Mexico City Cosmic Ray Observatory (cosmics ray fluxes), GPS receiver networks (ionospheric disturbances), and the Geomagnetic Observatory of Teoloyucan (geomagnetic field). The observations show that we detected significant space weather effects over the Mexican territory: geomagnetic and ionospheric disturbances (22 June, 2015), variations in cosmic rays fluxes, and also radio communication's interferences (29 September, 2015). The effects of these perturbations were registered, for the first time, using space weather products by SCiESMEX: TEC maps, regional geomagnetic index K$_{\rm mex}$, radio spectrographs of low frequency, and cosmic rays fluxes. These results prove the importance of monitoring space weather phenomena in the region and the need to strengthening the instrumentation network. | Space weather (SW) phenomena influence the performance and reliability of different modern technological systems; see for instance \cite{1999SSRv...88..563B} and \cite{Dinardini2016a}. The country has developed a significant infrastructure that is vulnerable to SW events, such as electricity generation and a transportation grid, telecommunications, electronic banking, long pipelines for gas and oil transportation, etc. The effects of the Carrington geomagnetic storm in 1859 were registered in several locations, indicating that the region is vulnerable to extreme geomagnetic storms \citep{gonzalez2018}. There are some studies of particular SW events that affected the geomagnetic field and ionosphere in Mexico, for example \cite{Rodriguez-Martinez2014}, \cite{2015AdSpR..55..586L}, \cite{SERGEEVA2017}, \cite{ROMEROHERNANDEZ2017}, and \cite{SERGEEVA2018}); however, the SW phenomena in this region have not been studied comprehensively. For instance, there is a lack of continuous multi-instrument observations of SW phenomena in Mexico that can provide reliable statistics for regional SW studies. Mexico is situated at low latitudes (geographic latitudes 14$^{\circ}$- 32$^{\circ}$N, geomagnetic latitudes 23$^{\circ}$ - 38$^{\circ}$ N). Recent studies prove that the SW effects are far from being fully understood at these latitudes \citep{2014JSWSC...4A..28C, SWE:SWE20063, balch2004halloween,balch2004intense}. The % southern half of Mexican territory is located between the Northern Tropic and the Equator. The Sun's incident ray path, at maximum elevation, remains troughout the year between 35$^{\circ}$ and 81$^{\circ}$ in the northern region of the country (Tijuana at 32$^{\circ}$N) and between 53$^{\circ}$ and 90$^{\circ}$ in the southern region (Tapachula at 14$^{\circ}$N). These conditions % match countries with % similar latitudes such as north of Africa, the Arabian Peninsula and south of Asia (including south of China and India). The Sun's paths for these latitudes increase the exposition time of the solar projection over the ground. Consequently, this raises the probability of radio interferences detected at ground level, produced directly or indirectly by a solar radio burst, Solar Energetical Particles (SEPs), and flares \citep{Lanzerotti2007}. Models like D-Region Absorption Predictions (D-RAP) by NOAA show this effect \citep{swpcRAP}. Since 2014, Mexico as begun a strategy for SW awareness. In 2014, the Mexican Space Weather Service (SCiESMEX) was created; in 2016, the National Space Weather Laboratory (LANCE) and the Repository of Space Weather Data (RICE) were established \citep{SWE:SWE20412}. Some of the ground-based instruments involved in the SW observations in Mexico have been used for more than 50 years \citep{Dinardini2016a,Dinardini2016b,Dinardini2016c}. Currently, the instrumental network provides the possibility of measuring local geomagnetic field variations, cosmic ray flux, solar wind parameters using interplanetary scintillation data, solar radio bursts, radio interferences, GPS signals delays, etc. The aim of this work is to estimate the impact of SW phenomena over % Mexico. We based our results on multilateral observations performed by the SCiESMEX instrumental network. In this work, we addressed two events registered over Mexico by SCiESMEX in 2015: on 22 June, and on 25-29 September. The first event was mainly related to a M6.5 solar flare and a geomagnetic storm that caused ionospheric perturbations. The second event was related to a solar radio burst. The paper is organized as follows: Section \ref{instrumentation} introduces the network of SW instruments in Mexico. Section 3 discusses the results of observations during the two events, and final remarks are given in the Conclusions section. | We presented the results of the first joint observations of SW phenomena in Mexico. We addressed two SW events that occurred on 22 June, 2015 and 29 September, 2015. Features of the behavior of SW parameters were obtained with the use of different local instruments installed in Mexico. The main results are the following: \begin{itemize} \item A solar flare was detected by the MEXART radio telescope on 22 June, 2015, in agreement with GOES satellite data. This example proves the possibility of using MEXART for solar flare detection if they occur during the local daylight hours. \item We presented for the first time a solar radio event (29 September, 2015) detected by the MEXART radio telescope and confirmed by the CALLISTO-MEXART station. The measuments by CALLISTO-MEXART were in accord with other CALLISTO observations. This proves that both ground-based local instruments (MEXART and CALLISTO-MEXART) can be used for the monitoring of solar radio bursts, which occur during the local daylight hours in Mexico. The advantage of the MEXART instrument is better sensitivity for such events. Note also that we report, for the first time, a radio blackout over Mexico related to SW phenomena. \item Local cosmic rays data indicate SW phenomena in Mexico. This is due to the fact that the irregularities in the interplanetary magnetic field, associated with large-scale solar wind disturbances, deflect the cosmic ray flux measured in the center of Mexico. For example, the record of Forbush decrease associated with the passing of the CME detected during the event in June 2015. \item Local geomagnetic field variations from 21 to 25 June, 2015 caused an intense ionospheric disturbance over Mexico. Local magnetometer data were in accord with the variations of the Dst index. The regional K$_{\rm mex}$ index allowed us to estimate the rate of geomagnetic disturbance in Mexico. The phases of ionospheric disturbance correlated with the phases of geomagnetic disturbance in time. The results are in agreement with other ionospheric studies of this event. It was revealed that the structure of the ionosphere was significantly changed during the geomagnetic storm which could lead to negative consequences for different technological systems. As the ionosphere state was estimated with only TEC data, no conclusions about the changes in each ionospheric layer can be made. \end{itemize} Some lessons can be learned from this first study in order to enhace the SW monitoring and the development of a comprehensive ground-based multi-instrument data set in Mexico. We must increase the number of magnetometers, located at different sites, to have local measurements at different regions in real time. The installation of more CALLISTO stations in Mexico will allow us to understand the effects of radio comunications disruption with more accuracy. For the case of TEC maps computed over Mexican territory, the next step is to improve the spatial resolution of TEC maps by increasing the number of GPS stations available and by bettering our TEC calibration methods. One of the future steps for improving the computations of TEC maps is to implement a homogeneous distribution of GPS stations throughout the ground territory. More detailed analysis could be made with the ionospheric sounding data by ionosondes. Thus, the ionosonde data are needed to complement both radio blackout studies and ionospheric radio propagation conditions over Mexico. The incorporation of a network of magnetometers and ionosondes in Mexico in the next year will significantly improve the coverage and quality of our space weather data. | 18 | 8 | 1808.07425 |
1808 | 1808.09324_arXiv.txt | { Meta-stable dark sector particles decaying into electrons or photons may non-trivially change the Hubble rate, lead to entropy injection into the thermal bath of Standard Model particles and may also photodisintegrate light nuclei formed in the early universe. We study generic constraints from Big Bang Nucleosynthesis on such a setup, with a particular emphasis on MeV-scale particles which are neither fully relativistic nor non-relativistic during all times relevant for Big Bang Nucleosynthesis. We apply our results to a simple model of self-interacting dark matter with a light scalar mediator. This setup turns out to be severely constrained by these considerations in combination with direct dark matter searches and will be fully tested with the next generation of low-threshold direct detection experiments.} | It is well known that fundamental particles leave their imprint on cosmological probes such as Big Bang Nucleosynthesis (BBN) or Cosmic Microwave Background (CMB) observations. In fact the remarkable agreement between the measured abundances of light elements such as deuterium and helium and the corresponding predictions within the Standard Model (SM) of particle physics implies that from about one second after the Big Bang the SM provides a good description of the early cosmological evolution. In particular, any deviations from the SM up to MeV-scale energies are strongly constrained~\cite{2016RvMP...88a5004C,Olive:2016xmw,Shvartsman:1969mm,Steigman:1977kc,Scherrer:1987rr,1988ApJ...331...33S}. There are various effects which influence the light element abundances. Of particular relevance is the expansion rate $H$ of the Universe during the time of BBN, as it determines at which point in time protons and neutrons fall out of thermodynamic equilibrium and hence sets the ratio of the corresponding number densities. The Hubble rate, in turn, is fully determined by the total energy density, which receives contributions from all particles, including those beyond the SM. In particular, even fully decoupled dark sectors can be probed via their effect on the expansion rate, a scenario which has been explored in detail recently~\cite{Hufnagel:2017dgo}. In this paper we extend this study to a different class of models, where additional effects from entropy production as well as the destruction of light nuclei via photodisintegration are relevant. Specifically we study the effect of exotic particles which are meta-stable and decay into electrons and photons during or after BBN. While previous studies have investigated similar scenarios~\cite{Sarkar:1984tt,Ellis:1984er,Scherrer:1987rr,Kawasaki:1994sc,Cyburt:2002uv,Jedamzik:2006xz,Poulin:2015woa,Poulin:2015opa}, it has always been assumed that the decaying particles are {\it non-relativistic} during BBN. Here we study the {\it fully general} case without such simplifying assumptions, which turns out to be relevant for a number of phenomenological applications. To this end, we develop in section~\ref{sec:cosmo_evolution} the formalism for the cosmological evolution of an MeV-scale particle $\phi$ which decays into $e^+ e^-$ and/or $\gamma \gamma$ with a lifetime $10^{-2} \, \text{s} \lesssim \tau_\phi \lesssim 10^8\,\text{s}$. To evaluate the effect on BBN, the calculation of the modified Hubble rate, as well as the non-standard evolution of the baryon-to-photon ratio $\eta$ originating from the production of entropy in the decays of $\phi$ are of particular relevance. In section~\ref{sec:abundance_of_light_elements} we evaluate the impact of the modified Hubble rate as well as the non-standard time-dependence of $\eta$ on the light nuclear abundances using a modified version of \texttt{AlterBBN}~\cite{Arbey:2011nf}. We then study the additional modifications of those nuclear abundances due to photodisintegration reactions induced by the decay products of $\phi$, which become relevant at $t \gtrsim 10^4\,$s. To this end, we first study the cascade of MeV-scale photons, electrons and positrons interacting with CMB photons as well as the background electrons and nuclei. In this context it is mandatory to properly solve the coupled Boltzmann equations, as the often utilised `universal photon spectrum' does not apply for the parameter values of interest as shown in~\cite{Poulin:2015woa, Poulin:2015opa}. We then compute the time evolution of the abundances of light elements which can be destroyed or created by collisions with those additional photons. Our general results are obtained by comparing the predicted nuclear abundances to observations for varying particle mass, abundance and lifetime in section~\ref{sec:results}. We also take into account the possibility of a different dark sector temperature and provide additional material in appendix~\ref{app:plot_collection} for convenience. In section~\ref{sec:SIDM} we apply our general results in conjunction with bounds from direct dark matter (DM) searches to a simple model of DM coupled to a light scalar mediator, which features potentially large DM self-interactions and has been extensively studied in the literature~\cite{Buckley:2009in,Loeb:2010gj,Kaplinghat:2013yxa,Kainulainen:2015sva,Kahlhoefer:2017umn}. Finally, in appendix~\ref{app:rates_cascade} we provide the rates for all processes relevant to the cascade of MeV-scale photons, electrons and positrons for reference, correcting a couple of typos found in the literature. | \label{sec:conclusions} Particles with a mass in the MeV range are predicted by various extensions of the Standard Model, with their coupling strength to SM states often being subject to stringent upper limits from direct or indirect searches. Unless one invokes additional even lighter (dark) states into which such a new particle $\phi$ can decay, this generically implies macroscopic lifetimes $\tau_\phi \gtrsim \mathcal{O}(1\,\text{s})$. While it is well known that this can potentially be in conflict with the remarkable success of standard Big Bang Nucleosynthesis, detailed studies of the impact of extra particles on the abundances of light nuclei so far have only been conducted in the limiting cases where the particle is either non-relativistic during BBN, or where it is ultra-relativistic and decays only well after BBN. However, when considering an MeV-scale particle, the four energy scales set by the particle mass $m_\phi$, the temperature during BBN, the temperature at the time of decay, as well as the binding energy of light nuclei such as deuterium can all be similar, which significantly complicates the physics underlying the calculation of BBN constraints. Motivated by this, we present for the first time a comprehensive study of BBN constraints on MeV-scale particles decaying into $e^+ e^-$ or $\gamma \gamma$ with a lifetime in the range $10^{-2}\,\text{s} < \tau_\phi < 10^8\,\text{s}$. To this end, we first numerically solve the full Boltzmann equation for the phase-space distribution function of the decaying particle $\phi$, without invoking any ultra- or non-relativistic approximation. Besides the additional energy density of $\phi$, which contributes to the expansion rate of the Universe and thus modifies the primordial abundances, we also consider in detail the production of entropy via the thermalisation of the decay products of $\phi$. This process can lead to a non-standard time-dependence of the baryon-to-photon-ratio $\eta(t)$ during the time when it is most relevant for BBN. We take into account both the enhanced Hubble rate as well as the non-standard baryon-to-photon ratio by properly modifying the public code \textsc{AlterBBN}. We find that in general both effects modify the predicted nuclear abundances in a similar way. Finally, if $m_\phi \gtrsim 4\,$MeV and $\tau_\phi \gtrsim 10^4\,$s, photodisintegration of light nuclei after the end of BBN can substantially modify the nuclear abundances. As already noted in~\cite{Poulin:2015woa,Poulin:2015opa}, for decaying particles with a mass in the MeV range the usually adopted `universal spectrum' of photons originating from the cascade process on the background photons is typically not applicable. Thus, by fully tracking the cascade evolution of high-energetic photons, electrons and positrons via double photon pair creation, photon-photon scattering, Bethe-Heitler pair creation, Compton and inverse Compton scattering, we derive the non-thermal photon spectrum and the associated photodisintegration rates of deuterium and helium separately for each point in parameter space for a given branching ratio of $\phi$. We then derive model-independent upper bounds on the ratio of the initial abundance of the particle $\phi$ and the photon-number density $n_\phi/n_\gamma$. To this end, we employ recent data on primordial abundances, and take into account systematic uncertainties on the nuclear rates relevant to BBN. Depending on the region in parameter space either of the effects related to the increased Hubble rate, the modified baryon-to-photon ratio or the photodisintegration of light nuclei can dominate the final constraint, reinforcing the necessity of a dedicated study of BBN constraints on MeV-scale particles. Importantly, when fixing the abundance of $\phi$ to the value expected for a thermal relic, we find that our upper limits in large parts of the parameter space deviate significantly from the frequently adopted order-of-magnitude estimates $\tau_\phi \lesssim 1\,$s (corresponding to the start of BBN) or $\tau_\phi \lesssim 10^4\,$s (corresponding to the start of photodisintegration). In appendix~\ref{app:plot_collection} we provide upper limits on $n_\phi/n_\gamma$ for a large set of model parameters, enabling the reader to quickly read off the BBN upper bound on the abundance of an unstable particle decaying into $e^+ e^-$ or $\gamma \gamma$. Lastly, we apply our general results to a specific model of self-interacting dark matter involving a fermionic dark matter particle $\psi$ interacting with a scalar mediator $\phi$, with the latter also having Higgs-like couplings to SM states. Such a scenario is compelling as it leads to large self-interaction cross sections of dark matter on small scales (and thus potentially solves tensions found within the pure $\Lambda$CDM model), while being consistent with upper bounds on the scales of galaxy clusters. However, the coupling strength of $\phi$ to SM particles is strongly constrained by direct detection experiments, rare kaon decays and bounds from the duration of the neutrino pulse from SN1987A. Based on a careful calculation of the cosmological evolution of both the dark matter particle $\psi$ and the unstable mediator $\phi$, we then derive for the first time detailed BBN constraints on this scenario. Our results show that almost all of the parameter space of the model leading to significant self-interactions of dark matter is ruled out by the combination of direct detection experiments and BBN, with only a small region around $m_\psi \simeq 0.5\,$GeV, $m_\phi \simeq 1.1\,$MeV and $\tau_\phi \simeq 30\,$s remaining. Interestingly, this combination of parameters can be fully tested with upcoming low-threshold direct detection experiments such as the final phase of CRESST-III~\cite{Angloher:2015eza}.\\[0.5cm] \textit{Note added:}\\ Shortly after the completion of this work, constraints from photodisintegration arising from the decay of MeV-scale particles were also studied in~\cite{Forestell:2018txr}, including for the first time the effect of FSR of photons. While this does not affect our bounds for $\tau_\phi \lesssim 10^4\,$s or $m_\phi \lesssim 5\,$MeV arising from the increased Hubble rate or entropy production (which is not considered in~\cite{Forestell:2018txr}), it does have an impact on the limits for a sufficiently heavy particle decaying into $e^+ e^-$ with a lifetime $\tau_\phi \gg 10^4\,$s. In this updated version of our work, we have thus included FSR via eq.~(\ref{eq:SFSR}), leading to minor changes in the left panel of Fig.~\ref{fig:results_fixed_mphi}, the right panel of Fig.~\ref{fig:results_fixed_tauphi} and the lower left panel of Fig.~\ref{fig:appendix_plots}. | 18 | 8 | 1808.09324 |
1808 | 1808.07878_arXiv.txt | Hypervelocity stars (HVSs) represent a unique population of stars in the Galaxy reflecting properties of the whole Galactic potential. Determining their origin is of fundamental importance to constrain the shape and mass of the dark halo. The leading scenario for the ejection of HVSs is an encounter with the supermassive black hole in the Galactic Centre. However, new proper motions from the \textit{Gaia} mission indicate that only the fastest HVSs can be traced back to the Galactic centre and the remaining stars originate in the disc or halo. In this paper, we study HVSs generated by encounters of stellar binaries with an intermediate-mass black hole (IMBH) in the core of a star cluster. For the first time, we model the effect of the cluster orbit in the Galactic potential on the observable properties of the ejected population. HVSs generated by this mechanism do not travel on radial orbits consistent with a Galactic centre origin, but rather point back to their parent cluster, thus providing observational evidence for the presence of an IMBH. We also model the ejection of high-velocity stars from the Galactic population of globular clusters, assuming that they all contain an IMBH, including the effects of the cluster's orbit and propagation of the star in the Galactic potential up to detection. We find that high-velocity stars ejected by IMBHs have distinctive distributions in velocity, Galactocentric distance and Galactic latitude, which can be used to distinguish them from runaway stars and stars ejected from the Galactic Centre. | The European Space Agency mission \textit{Gaia}\footnote{http://sci.esa.int/gaia/} has revolutionized astrometry, providing positions, parallaxes and proper motions for more than $1.3$ billion stars in its second data release (Gaia DR2) \citep{gaia18}. It also provided radial velocities for $\sim 7.2$ million bright stars \citep{gaia18}. It therefore offers an unprecedented opportunity to study the population of high-velocity stars in the Galaxy. Based on their ejection mechanism and/or space velocities, high-velocity stars are usually divided in two different categories: runaway stars (RSs) and hypervelocity stars (HVSs), though the distinction is not always clear. The former likely originated in the Galactic disc and acquired high velocities either in supernova explosions in binary systems \citep{port00} or in close dynamical encounters involving stars and binaries \citep{gvar09,gvar11}. The latter, on the other hand, have such extreme velocities, often exceeding the local Galactic escape speed, that an interaction with the supermassive black hole (SMBH) in the Galactic Centre (GC) is required. The leading scenario, the classical Hills mechanism \citep{hills88}, involves the tidal disruption of a stellar binary by the SMBH which results in the capture of one star on a wide eccentric orbit and the ejection of the companion with very large velocity \citep{brm06,sari10}. Alternative mechanisms for the origin of HVSs include encounters with a massive black hole binary in the Galactic Centre \citep{yut03,fraglei18}, inspiralling intermediate-mass black holes \citep[IMBHs;][]{baum2006,sesa2006,sesan07}, encounters in a nearby galaxy \citep{GPZ2007,she2008,bou2017,erk19} and tidal interactions of stars clusters with a single or binary SMBHs \citep{cap15,fra16,fck17}. Unbound stars originating from the Galactic disc have also been discovered, e.g. HD 271791 \citep{heber08}, and these are often named hyperrunaway stars (HRSs). More than $20$ early-type HVSs have been confirmed by the spectroscopic Multiple Mirror Telescope (MMT) survey, with Galactocentric velocities up to $\sim 700\kms$ and distances between $\sim 50$ and $\sim 120\kpc$ from the GC \citep{brw06,brw12,brw14}. Recently, \citet{boub18} used \textit{Gaia} data to show that late-type HVS candidates \citep[see e.g.][]{silva11,zhong14,vickers15} are likely bound to the Milky Way, except for LAMOST J115209.12+120258.0 \citep{li15}, which moves on an unbound orbit not originating in the GC. Determining the exact origin of HVSs is of extreme importance for studies of the Galactic mass distribution and dark halo. If originating in the GC and travelling on almost radial orbits to the halo, HVSs can be used to probe the shape of the Galactic potential \citep{gnedin2010,fl2017,rossi2017}. \citet{brown18} show that only the fastest HVSs (with radial velocities $\gtrsim 450\kms$) have orbits originating in the GC, while the other unbound stars in the MMT sample have ambiguous origin. \citet{march18} used \textit{Gaia} data to identify HVS-candidates and found $28$ objects out of $165$ with a significant ($\gtrsim 50\%$) probability of being unbound, of which $\sim 2-5$ come from the GC and $\sim 9$ from the Galactic disc, and the remaining are likely of extragalactic origin. While HVSs coming from other galaxies and satellites (as the Large Magellanic Cloud; LMC) may contaminate the sample \citep{erk19}, \citet{keny18} showed that the Galactic disc and the LMC potential may have a role in deflecting HVSs from a nearly radial orbit. Yet, the origin of most of the known HVSs remains unknown \citep{brown18}. A possible origin for the Galactic HVSs that cannot be traced back to the GC may be a Hills-type process in star clusters hosting an IMBH. Assuming that the observed $M_{SMBH}-\sigma$ relation (with $\sigma$ the local stellar velocity dispersion) holds also for the range of IMBH masses ($10^2\msun \lesssim M_{IMBH} \lesssim 10^5 \msun$), star clusters are the place where IMBHs should reside \citep*{fragk18,frlgk18}. The recent observation of a tidal disruption event consistent with an IMBH in an off-centre star cluster \citep{lin18} represents a significant milestone in the hunt for the elusive IMBH population. In these clusters, binaries are disrupted by the tidal field of the IMBH and are ejected from the cluster with high velocities. In this paper, we study the ejection of HVSs from encounters between binary stars and an IMBH, as first proposed by \citet{pfa05} and \citet{GPZ2007}, in the core of star clusters by means of high-precision scattering experiments. Using a similar approach, \citet{sesan12} suggested that IMBH in star clusters may eject a population of millisecond pulsars to the Galactic halo. We study the imprint of the binary mass, binary semi-major axis and IMBH mass on the spatial and velocity distributions of the stars ejected from the cluster. Moreover, for the first time, we consider the effect of the cluster orbit through the Galaxy on the observable properties of the high-velocity stars, which can contribute both to the Galactic RS and HVS populations. We show that HVSs generated by a Hills-like mechanism in a star cluster hosting an IMBH would not travel on radial orbits consistent with a Galactic Centre origin, rather they would point back to their parent cluster, thus revealing the presence of an IMBH. We also show that high-velocity stars ejected by IMBHs have distinctive distributions in velocity, Galactocentric distance and Galactic latitude, which can be used to distinguish them from high-velocity stars from other channels. The paper is organised as follows. In Section \ref{sect:hills}, we describe the classical Hills model for binary breakups by an IMBH. In Section \ref{sect:numsetup}, we describe the methods used in our numerical experiments. In Section \ref{sect:distributions}, we present the velocity distributions of HVSs after the tidal breakup of a binary by an IMBH, while, in Section \ref{sect:clusters}, we discuss the impact of the host cluster orbit on the HVS spatial distribution. In Section \ref{sect:comparisons}, we compare the properties of the high-velocity stars produced by the Galactic population of globular clusters to the high-velocity stars originated both in the GC and the Galactic disc. Finally, a discussion and conclusions are presented in Section \ref{sect:concl}. | \label{sect:concl} The \textit{Gaia} mission has revolutionized astrometry and is offering an unprecedented opportunity to study the population of high-velocity stars in our Galaxy, thanks to its precise proper motion measurements. The full \textit{Gaia} data release will allow to determine the origin of (most) HVSs and put constraints on the Galactic mass distribution and dark halo potential \citep{gnedin2010,fl2017,rossi2017}. The leading scenario to explain the velocities of the hypervelocity stars in the halo is the classical Hills mechanism \citep{hills88}, where a stellar binary is unbound by the SMBH in the Galactic centre, as a result of which one star is captured in orbit around the SMBH while the companion is ejected with a typical velocity of a thousand $\kms$. This model predicts HVSs moving away from the Galactic centre on nearly radial orbits. Surprisingly, recent observations show that only a few of the known HVSs (the ones with the highest velocities) can be traced back to the centre of the Galaxy \citep{boub18,brown18,march18, erk19}. \citet{keny18} showed that the Galactic disc and the LMC potential may play a role in deflecting HVSs from a nearly radial orbit. Yet, the origin of most of the known HVSs remains unknown. We considered the ejection of high velocity stars from star clusters hosting an IMBH due to encounters between binary stars and the IMBH itself. We performed a large number of high-precision scattering experiments varying the IMBH mass and the properties of the binaries and derived the velocity and space distributions of the ejected stars by taking into account the effect of the parent cluster's orbital motion. We found that the properties of the ejected stars, which can contribute both to the Galactic RS and HVS populations, depend on the cluster orbit, with most of the unbound HVSs found at Galactocentric distances $R\gtrsim 30$ kpc. As expected, HVSs generated by a Hills-like mechanism in a star cluster hosting an IMBH would not travel on orbits consistent with a Galactic centre origin, rather they would point back to their original host cluster, thus providing observational evidence for the presence of an IMBH. We caution, however, that observational errors in the distance, radial velocity and proper motions would propagate while backtracing the orbit in the Galactic potential, thus making a precise localisation of the birth location in the disc challenging. In addition to a population of unbound HVSs, this mechanism can produce a large population of bound RSs \citep*{silva11,zhong14,vickers15,subr2019}. We also modelled the ejection of high-velocity stars from the Galactic population of globular clusters, assuming that each cluster contains an IMBH and combining the ejection velocity of the star with the cluster's orbital velocity, as well as propagation of the star in the Galactic potential up to the estimated detection time. We compared the properties of the resulting high-velocity stars with those of mock populations ejected from both the Galactic disc (dynamical and SN ejections) and the Galactic centre (Hills mechanism). We find that high-velocity stars ejected by IMBHs have distinctive distributions in velocity, Galactocentric distance and Galactic latitude. The ejection rate is $\sim 0.7-1.2\times10^{-4}$ yr$^{-1}$, similar or slightly larger than the rate for the Hills mechanism. Altogether, this opens up the possibility of constraining the origin of high-velocity stars with future observations. A detailed comparison with available data requires a detailed analysis of observational selections and biases and is left to future work. Studying spectroscopic and kinematic data of runaways and HVSs can help constraining their origin and the possible presence of an IMBH in the core of a given globular cluster. Spectroscopic data can be used to obtain radial velocities, while tangential velocities can be obtained from proper motion measurements, whose combination, along with the sky position of the stars, gives the full 6-D phase space information to determine their trajectories. This approach has already proven successful in a few cases in which high-velocity stars were traced back to their birth location, providing indirect evidence for the existence of IMBHs in certain clusters. \citet{hoo01} used radio observations and milli-arcsecond accuracy astrometry from the \textit{Hipparcos} satellite of the orbits of $56$ RSs and nine compact objects with distances $\lesssim 700$ pc to identify their parent stellar group. \citet{hebe08} studied the mass, evolutionary lifetime and kinematics of HD 271791 by using proper motion measurements from a collection of catalogues, and found that the likely birthplace of HD 271791 is the outer Galactic disc, while the Galactic Centre origin is ruled out. Recently, \citet{lenno18} and \citet{renz18} used \textit{Gaia} data to show that the dynamics of the very-massive runaways VFTS 16 and VFTS682 are consistent with ejection from the young massive cluster R136. \citet{hatt18} studied the origin of hyper-runaway subgiant LAMOST-HVS1 by using \textit{Gaia} data, and found that it was likely ejected dynamically from near the Norma spiral arm as a consequence of a few-body encounter or a Hills-like ejection by an IMBH. Thanks to present and future high precision data from \textit{Gaia}, similar studies can be performed for known and candidate high-velocity stars, which might have been ejected through the mechanism discussed in this work, thus revealing new IMBH candidates. | 18 | 8 | 1808.07878 |
1808 | 1808.02493_arXiv.txt | The conditions and evolution of protoplanetary disks in multiple systems can be considerably different from those around single stars, which may have important consequences for planet formation. We present Very Large Array (VLA) 8.8\,mm (34\,GHz) and 5\,cm (6\,GHz) observations of the quadruple system HD~98800, which consists of two spectroscopic binary systems (Aa-Ab, Ba-Bb). The Ba-Bb pair is surrounded by a circumbinary disk, which is usually assumed to be a debris disk given its $\sim$10\,Myr age and the lack of near infrared excess. The VLA 8.8\,mm observations resolve the disk size (5-5.5\,au) and its inner cavity ($\approx$3\,au) for the first time, making it one of the smallest disks known. Its small size, large fractional luminosity, and millimeter spectral index consistent with blackbody emission support the idea that HD~98800~B is a massive, optically thick ring that may still retain significant amounts of gas. The disk detection at 5\,cm is compatible with free-free emission from photoionized material. The diskless HD~98800~A component is also detected, showing partial polarization at 5\,cm that is compatible with nonthermal chromospheric activity. We propose that tidal torques from Ba-Bb and A-B have stopped the viscous evolution of the inner and outer disk radii, and the disk is evolving via mass loss through photoevaporative winds. This scenario can explain the properties and longevity of HD~98800~B as well as the lack of a disk around HD~98800~A, suggesting that planet formation could have more time to proceed in multiple systems than around single stars in certain system configurations. | \label{sec:intro} Given the abundance of binaries and multiple systems in our Galaxy, understanding planet formation in these systems represents an important piece of our knowledge of exoplanetary populations. In fact, 30-75\,\% of stars appear to have formed in multiple systems \citep{Duquennoy1991, Kraus2008, Lafreniere2008, Kraus2011, Duchene2013}, and various exoplanets have already been found around binary and multiple systems \citep[e.g.][]{Doyle2011, Welsh2012_kepler_binaries, Dupuy2016, Kostov2016}. The formation of planets around single stars occurs in the protoplanetary disks surrounding young stars. These disks contain gas and dust which, via dust growth first and gas accretion later, can form rocky and gas giant planets \citep{Raymond2014, Helled2014, Testi2014}. At the same time, processes such as the viscous evolution of the disk \citep[e.g.][]{Shakura1973, Lynden-Bell1974, Hartmann1998}, photoevaporation by stellar radiation \citep[e.g.][]{Shu1993, Hollenbach1994}, or the interaction with newborn planets \citep[][]{Kley2012, Espaillat2014} lead to the eventual dispersal of the disk \citep[see][and references therein]{Alexander2014}, which occurs 5-10\,Myr after their formation \citep{Haisch2001, Hernandez2007, Hernandez2008, Mamajek2009, Ribas2015}. Once the protoplanetary disk disperses, a second generation of dust produced by planetesimal collisions (with some possible contribution from remaining dust from the protoplanetary phase) forms a gas-poor, less massive, colder disk called a debris disk, analogous to our own Kuiper Belt \citep[e.g.][]{Wyatt2008, Hughes2018}. While the latter can survive for much longer ($\sim$Gyr) thanks to dust replenishment from these collisions, the transition from a protoplanetary to a debris disk is still not well understood. The importance of disk evolution and dispersal is obvious, since it sets a time limit to planet formation and determines many of the properties of the resulting planetary systems. This scenario can change significantly in binary or multiple systems; the gravitational interaction of a (sub)stellar companion can have an important effect on disks, leading to truncation or even quick dispersal depending on various parameters, such as the semi-major axis of the binary orbit, its eccentricity, or the masses of the components \citep[e.g.][]{Papaloizou1977, Artymowicz1994}. In fact, various surveys have shown the fraction of protoplanetary disks to be smaller around close ($a< 40$\,au) binaries than around single stars or wide binaries \citep{Cieza2009, Harris2012, Kraus2012}. On the other hand, very close binaries (a few au) may not be able to retain individual circumstellar disks (disks around individual components of the system), but could harbor a circumbinary disk surrounding both stars. Mass accretion rates in circumbinary disks can be severely diminished due to the tidal torques exerted by the central sources, resulting in extended disk lifetimes \citep[e.g.][]{Alexander2012}. Also, disks with truncated outer radii due to a close companion may even experience an outside-in evolution \citep{Rosotti2018}, which is completely different from the expectation for disks around single stars. All these mechanisms are likely to result in different populations of exoplanets around binary and multiple stars and are thus key pieces for understanding them. In this study, we present NSF's Karl G. Jansky Very Large Array (VLA) observations of an extreme case of a circumstellar disk in a multiple system: the hierarchical quadruple system HD~98800. The VLA observations resolved both the size and the inner cavity of the disk around HD~98800~B for the first time and shed light onto the puzzling existence of long-lived disks in multiple systems. The paper is outlined as follows: Sec.~\ref{sec:description} provides a description of the HD~98800 system. Sec.~\ref{sec:observations} describes the VLA observations and ancillary data used, and the results and modeling procedure are discussed in Sec.~\ref{sec:results}. In Sec.~\ref{sec:discussion}, we examine the nature of the disk around HD~98800~B, compare our results with previous studies, and discuss implications for disks in multiple systems. Finally, our conclusions are presented in Sec.~\ref{sec:conclusions}. | \label{sec:conclusions} We present Ka~band (8.8\,mm, 34\,GHz) and C~band (5\,cm, 6\,GHz) VLA observations of the hierarchical quadruple system HD~98800 in the 10\,Myr old TW~Hya association. The 8.8\,mm observations successfully resolve the disk and the inner cavity around the B component and detect unresolved emission from A. Both A and B are also detected (unresolved) at 5\,cm. Only the 5\,cm emission from A shows signatures of polarization. We use the MCFOST radiative transfer software to model both the SED and the resolved VLA image of the disk. The main results of the study are: \begin{enumerate} \item The (sub)mm spectral index of the disk around HD~98800~B is $\approx2$ up to 8.8\,mm, suggesting either optically thin radiation from a dust population with a gray opacity law at millimeter wavelengths or optically thick emission. \item The disk extends from 3 to 5\,au at 8.8\,mm. The values of both the inner and outer radii (truncated by the Ba-Bb and the interaction with A, respectively) agree with model predictions of viscous disk interactions with companions. \item For such a disk size, we were unable to reproduce the observed 8.8\,mm flux and spectral index with an optically thin disk. Combined with previous detections of gas in the system, this suggests that HD~98800~B is probably a massive, optically thick and gas-rich disk, more similar to protoplanetary than to debris disks. \item Given that both disk fractions and exoplanet occurrence are lower in close binaries than around single stars, the survival of a massive disk in a quadruple system, such as HD~98800~B, appears to contradict disk evolution theories. We suggest that the inner and outer truncation of the disk have stopped the viscous evolution of the disk, which is thus governed by photoevaporation. The small size of the disk lowers the mass-loss rate through photoevaporation, explaining the longevity of the disk. This scenario results in longer disk lifetimes that may ease planet formation in multiple systems with the appropriate configuration. \item The spectral index of HD~98800~B between 8.8\,mm and 5\,cm is significantly shallower ($\alpha=1.3$), indicating that additional contribution from free-free emission by photoevaporating material or from gyrosynchrotron emission from stellar activity is present at cm wavelengths. In case the photoevaporative wind is the correct explanation, then it is possible that the 8.8\,mm flux is also partially probing this emission, making HD~98800~B an ideal target to study this phenomenon. \item The disk size is 2-3\,times smaller than the one derived by \citet{Andrews2010a}, who used partially resolved SMA observations at 880\,$\mu$m. While we attribute the difference to the lower angular resolution of SMA, if this size dependence with wavelength is, in fact, real, it may hint at significant dust radial migration in the system. \item The A component appears unresolved both at 8.8\,mm and 5\,cm, and the observed fluxes are above the expected photospheric levels. In contrast to HD~98800~B, the spectral index between the two VLA bands is negative ($\alpha=-0.4$). Together with the partial polarization of the 5\,cm emission of HD~98800~A, this favors a stellar origin for the excess instead of a previously unknown disk around this component. \end{enumerate} Follow-up ALMA observations will be able to search for CO in the system with increased sensitivity with respect to previous studies and to measure the disk size at shorter wavelengths to test the dust radial migration scenario in HD~98800~B. Additional VLA observations at other wavelengths will also provide a better sampling of the mm/cm range and will help to quantify the relative importance (if any) of free-free emission in the disk. Given its proximity and unique properties, HD~98800 offers a unique laboratory to study the impact of multiplicity on planet formation and disk evolution. \vspace{1cm} We thank the anonymous referee for their useful comments, which helped improving the quality of this manuscript. We thank Richard Alexander for discussion on disk evolution in binary systems and Eric Nielsen for examining the orbit of HD~98800 including the VLA data presented here. This manuscript is based on data from the NSF's Karl G. Jansky Very Large Array (VLA). We have also used data from the \emph{Herschel Space Observatory} (\emph{Herschel}). \emph{Herschel} is an ESA space observatory with science instruments provided by European-led Principal Investigator consortia and with important participation from NASA. \software{emcee \citep{emcee}, MCFOST \citep{MCFOST}, HIPE \citep[v15;][]{Ott2010}, CASA \citep{CASA}, Matplotlib \citep{Matplotlib}, SciPy \citep{Scipy}, Numpy \citep{Scipy}, pandas \citep{pandas}, Astropy \citep{Astropy2013}.} | 18 | 8 | 1808.02493 |
1808 | 1808.03389_arXiv.txt | The recent discovery of hot dust grains in the vicinity of main-sequence stars has become a hot issue among the scientific community of debris disks. Hot grains must have been enormously accumulated near their sublimation zones, but it is a mystery how such a high concentration of hot grains is sustained. The most difficult conundrum is that the size of hot dust grains is estimated to lie in the submicrometer range, while submicrometer-sized grains are instantly swept away from near-stellar environments by stellar radiation pressure. One and only mechanism proposed for prolonging the residence time of hot grains in the near-stellar environments is trapping of charged nanoparticles by stellar magnetic fields. We revisit the model of magnetic grain trapping around main-sequence stars of various spectral classes by taking into account sublimation and electric charging of the grains. The model of magnetic grain trapping predicts that hot dust grains are present in the vicinity of main-sequence stars with high rotation velocities and intermediate magnetic-field strengths. On the contrary, we find that the detection of hot dust grains has no correlation with the rotation velocities of central stars nor the magnetic field strengths of the stars. Our numerical evaluation of electric grain charging indicates that the surface potential of submicrometer-sized grains in the vicinity of main-sequence stars is typically $4$--$5~\mathrm{V}$, which is one order of magnitude smaller than the value assumed by the model of magnetic grain trapping. On the basis of our numerical simulation on sublimation of dust grains in the vicinity of a star, it turns out that their lives end due to sublimation in a timescale much shorter than the period of one revolution at the gyroradius. It is, therefore, infeasible to dynamically extend the dwell time of hot grains inside the sublimation zone by magnetic trapping, while we cannot completely rule out the possibility of magnetic grain trapping outside the sublimation zone where the strength of stellar magnetic field has been underestimated in the previous model. Nevertheless, the independence of hot dust detection on the stellar rotational velocity and magnetic field strength favors a scenario that some other (yet unnoticed/overlooked) ubiquitous mechanism of grain trapping is at work. | \label{sec:intro} The central region of debris disks around main-sequence stars is best visible by near-infrared wavelengths at which thermal emission from hot dust grains near the central stars peaks. Near-infrared interferometric observations of debris disks, in other words, exozodis\footnote{Here we use the term ``exozodis'' as an analog to the zodiacal dust cloud that extends from the Kuiper belt dust ring(s) to the solar F-corona.} have made the discovery of excess emission from hot dust in the vicinity of main-sequence stars \citep{absil-et-al2009,absil-et-al2013,ertel-et-al2014}. It turned out that the excess emission from hot dust equivalent to approximately 1\% of stelar flux is a serious challenge for classical grain dynamics models. Currently available parametric models of exozodis attribute the near-infrared excess to an enormous amount of submicrometer-sized grains near the sublimation zone, although they are too small to stay in bound orbits around the central stars against strong stellar radiation pressure \citep{lebreton-et-al2013,vanlieshout-et-al2014,kirchschlager-et-al2017}. In a classical grain dynamics model, dust grains in orbit around a central star gradually spiral to the star by the Poynting-Robertson effect and accumulate at the outer edge of sublimation zone on account of the increase in the semimajor axis by sublimation \citep{burns-et-al1979,kimura-et-al1997,kobayashi-et-al2009}. However, the enhancement of hot dust density at the outer edge of sublimation zone is limited to a factor of 10, which is far from sufficient to explain the observed excess emission from hot dust \citep{kobayashi-et-al2009,kobayashi-et-al2011,vanlieshout-et-al2014}. There are proposed mechanisms to supply dust grains in situ by electrostatic ejection of the grains from their parent bodies or through cometary activity of planetesimals in inner mean-motion resonances with eccentric planets \citep{kimura-et-al2014,faramaz-et-al2016}. Unfortunately, a vast supply of submicrometer-sized grains into the near-stellar environments would result in a considerable amount of dust outflows, which has not been observed to date. Moreover, these mechanisms are unable to prolong the resident time of the grains in the near-stellar environments, while trapping of dust grains in the vicinity of stars by a yet another mechanism, if exists, is relatively straightforward. \citet{rieke-et-al2016} have proposed a model of magnetic trapping in which charged nanoparticles are dynamically confined in the stellar magnetic field inside the sublimation zone. Nonetheless, we find it difficult to justify the underlying assumptions of the magnetic grain trapping; First, it is not obvious whether the production of nanoparticles in the vicinity of a star by whatever mechanisms results in magnetic trapping instead of radiative blowout. Second, the model explicitly assumes that dust grains around A-type stars obtain the same electrical charges as around young stellar objects (YSOs), irrespective of a significant difference in their plasma temperatures, on which the electrostatic potential may be dependent \citep[cf.][]{telleschi-et-al2005,draine-salpeter1979}. Third, the model implicitly assumes that the radiation spectrum of A-type stars follows Planck's law with an effective temperature of 10,000~K, while the deviation of stellar spectra from Planck's law is significant in the ultraviolet, which might affect the photoelectric charging. Fourth, the model ignores the effect of sublimation on the dynamics of hot dust grains, while sublimation may prevent the grains from being trapped for a considerable period of time. Fifth, the model suggests that the duration of magnetic grain trapping is on the order of months or a few years at most, but the resident time of hot nanoparticles expected from observations is 5--50 years \citep{su-et-al2016}. Sixth, the model predicts an association between the detection of near-infrared excess emission and the rotation velocity of stars, but rapid rotators do not always have a near-infrared excess \citep{absil-et-al2013}. Last but not least, since \citet{rieke-et-al2016} ignored collisional destruction and stellar-wind velocity in their simulations, the magnetic trapping scenario needs to be further investigated with detailed numerical simulations, as emphasized by \citet{kral-et-al2017}. Inevitably, dust particles in debris disks are electrically charged by virtue of exposure to ultraviolet radiation and plasma flow from central stars, irrespective of their physical and chemical properties. As a result, the magnetic field of the central star might influence the dynamics of dust particles by exerting a force, called the Lorentz force, on the particles, the smallest ones, in particular \citep{parker1964,belton1966}. This is because the Lorentz force is proportional to the radius of dust particles and the gravitational force is proportional to the volume of the particles. Numerical simulations on the dynamics of charged dust particles in the close vicinity of the Sun have shown that the Lorentz force plays an important role in the dynamics of the particles near the sublimation zone \citep{rusk1988,krivov-et-al1998,czechowski-mann2010,czechowski-kleimann2017}. The magnitude of the Lorentz force acting on a charged grain is proportional to the electrical charges on the surface of the grain or the electrostatic surface potential in other words. Therefore, it is clear that the electrical charge on the grain surface is the key quantity to elucidate the dynamics of dust particles near the sublimation zone in a stellar magnetic field. The electrostatic potential of dust particles, which is mainly determined by the current balance between plasma electrons and photoelectrons, has been estimated for the zodiacal dust disk around the Sun \citep{mukai1981,kimura-mann1998a}. In the vicinity of stars, thermionic emission and secondary electron emission might also make major contributions to the electric charging of grains \citep{belton1966,lefevre1975,mukai1981}. Unfortunately, it is not obvious how the electrostatic potential varies in the sublimation zone nor how the spectral class of main-sequence stars affects the potential. Accordingly, a proper estimate of electrical charges on the surface of dust grains could help us understand the dynamics of dust grains in the near-stellar environments. In this paper, we revisit the magnetic trapping of hot dust grains around main-sequence stars of various spectral classes by taking into account sublimation and electric charging of the grains. We first model electrostatic grain charging in the vicinity of main-sequence stars with various spectral types and stellar wind parameters. Subsequently we present our numerical results on the equilibrium temperatures, electrostatic potentials, and charge-to-mass ratios of dust particles near the sublimation zone, as well as the ratio of forces acting on the particles in the stellar magnetic and gravitational fields. The results are utilized to compare characteristic timescales for sublimation to periods of gyrating motion, and the detections of near-infrared excess emission to stellar rotation velocities and magnetic field strengths. Finally we discuss the plausibility of magnetic grain trapping based on our results of grain dynamics and our findings about the ubiquitous nature of hot grains. | We have revisited the magnetic grain trapping by modeling the electric charging of dust grains in the vicinity of main-sequence stars as the electric charge on the grain surface is the key parameter for the strength of magnetic grain trapping. The model utilizes our estimates of stellar parameters and stellar-wind parameters, which are requisites for computing electric currents of various charging processes. The most advantage of our model is that the computation of grain charges is not restricted to spectral classification of the stars nor physical properties of the grains. Therefore, one can apply our model of grain charging to any stellar environments and grain properties, as far as main-sequence stars are concerned. It has turned out that dust grains in exozodis are electrically charged up to $U \approx 4$--$5~\mathrm{V}$ at least for stellar environments and grain properties considered in this paper. Although nanometer-sized grains may attain higher surface potentials in the sublimation zone, we expect that the intense electrical charging has only a limited effect on the dynamics of the grains compared to sublimation. A typical value of $U \approx 4$--$5~\mathrm{V}$ in exozodis agrees with the surface potential of dust grains in the zodiacal cloud around the Sun \citep{mukai1981,kimura-mann1998a}. Consequently, if only an order of magnitude estimate is required for the surface potential of grains in exozodis, it is safe to assume the surface potential of $U \approx 4$--$5~\mathrm{V}$. Contrary to the electrostatic surface potential of $U = 40~\mathrm{V}$ assumed in the model of magnetic grain trapping by \citet{rieke-et-al2016}, we have revealed in Fig.~\ref{fig:2} that a typical value of surface potential be $U \approx 4$--$5~\mathrm{V}$ in exozodis. This stems from the fact that the major charging process is the photoelectron emission where the maximum energy of photoelectrons determines the surface potential. As a result, the model of magnetic grain trapping clearly overestimates the surface potential of nanoparticles in the sublimation zone by one order of magnitude. In contrast, they assumed the strength of magnetic field at stellar surface to be $B_\star = {10}^{-4}~\mathrm{T}$, while realistic strengths of stellar magnetic fields might be one to two orders of magnitude higher than the assumed value. Because the product of grain electric charge and magnetic field strength is the key quantity for the dynamics of hot grains in a stellar magnetic field, the motion of hot grains might be reasonably well simulated by the model of \citet{rieke-et-al2016}. Therefore, we cannot dismiss the model of magnetic grain trapping only by our numerical estimates of electric grain charging. We have shown that the sublimation of grains takes place in a much shorter timescale than the magnetic trapping of the grains inside the sublimation zone at least in the cases of $\beta$~Pic and $\epsilon$~Eri. The strength of magnetic field at the surface of $\beta$~Pic has been estimated to be $B_\star \approx 8 \times {10}^{-3}~\mathrm{T}$, more than one order of magnitude larger than the value assumed by \citet{rieke-et-al2016}. Therefore, the correction of the stellar magnetic field strength to the model of \citet{rieke-et-al2016} does not significantly reduce the period of gyration short enough to prolong the lifetime of grains trapped in the stellar magnetic field. In consequence, the magnetic trapping of grains in the sublimation zone does not seem to be a potential mechanism to prolong the lifetimes of the grains in the near-stellar environments. While we have assumed $B \propto r^{-3}$ according to \citet{rieke-et-al2016}, we are aware that the radial dependence of the magnetic field strength becomes weaker as the radial distance $r$ increases. Such a variation in the radial dependence of the magnetic field strength is associated with the nature of magnetic field lines frozen in the stellar wind. As a result, the azimuthal component of magnetic fields could predominate over the radial component in the far distance from the central star, which leads to $B \propto r^{-1}$. This might imply that the periods for gyration motion, $\tau_\mathrm{g}$, shown in Fig.~\ref{fig:4} be significantly overestimated in the far distance. In particular, the radial distance beyond which the azimuthal magnetic field becomes the dominant component is closer to the central star for fast rotating stars in comparison with slow rotators \citep{weber-davis1967,belcher-macgregor1976,sakurai1985}. It should be, however, noted that the assumption of $\Theta = 90\deg$ (i.e., $\mathbf{v_\mathrm{g}} \perp \mathbf{B}$) is out of harmony with the predominance of the azimuthal component. According to \citet{parker1958}, we may assume $\sin\Theta \approx (v_\mathrm{sw}^\infty/V_{\star}) (R_{\star}/r)$ in the far distance from the central star. This results in $\tau_\mathrm{g} \propto r^2$ at large $r$, instead of $\tau_\mathrm{g} \propto r^3$, indicating that we might have overestimated $\tau_\mathrm{g}$ by an order of magnitude, but not more than two orders of magnitude within $r = 1~\mathrm{au}$. Therefore, we still obtain $\tau_\mathrm{s} < \tau_\mathrm{g}$ inside the sublimation zone and thus rapid sublimation prevents the magnetic trapping of charged nanoparticles from taking place in this region. This does not exclude the possibility of magnetic grain trapping {\it outside} the sublimation zone where the strength of stellar magnetic fields has been underestimated in the model of \citet{rieke-et-al2016}. Even if hot dust grains accumulate outside the sublimation zone, the near-infrared excess emission from hot grains implies that the location of the grains should not be very far from the outer edge of the sublimation zone. In case that the Lorentz force acting on dust grains is smaller the gravitational force on the particles at distances very far from the outer edge of the sublimation zone, they must first be transported to near-stellar environments by the Poynting-Robertson effect. However, small dust grains cannot stay in bound orbits after released from their parent bodies, provided that the radiation pressure force exceeds half the gravitational force. By studying the ratio of forces acting on nanoparticles in the magnetic and gravitational fields of a central star, we confirm that the Lorentz force dominates the dynamics of nanoparticles smaller than $a=100~\mathrm{nm}$ near the outer edge of the sublimation zone (see Fig.~\ref{fig:3}). Therefore, the dynamics of nanoparticles, if they exist, is controlled by the magnetic field and only sublimation poses a threat to the scenario for the magnetic grain trapping. If the mechanism of trapping hot grains in the vicinity of a central star is not associated with stellar magnetic fields, then we must anticipate that the size of hot dust grains lies in the range of $a \ga 100~\mathrm{nm}$. One of our important results is that the observed interferometric data do not show any correlation between the detection of hot grains and the magnitude of stellar rotations nor the strength of stellar magnetic fields (Figs.~\ref{fig:6} and \ref{fig:7}). This is clearly at odds with the proposed model of magnetic grain trapping by \citet{rieke-et-al2016} that predicts the presence of hot grains around rapidly rotating stars with intermediate magnetic fields. It is worth noting that \citet{absil-et-al2013} suggested a possible correlation between the near-infrared excess emission and the rotation velocities of A stars, while \citet{ertel-et-al2014} claimed no correlation. The arguments in the latter are based on the presence of fast rotators without hot grains, but a temporal variation in the detection of hot grains may weaken their arguments. In contrast, our results reveal the presence of slow rotators with hot grains providing strong evidence for no correlation between the near-infrared excess emission and fast stellar rotation. Our results imply that the proposed model of magnetic grain trapping needs to be rejected or at least largely modified to get rid of the dependence on the stellar rotation and magnetic field strength. Since the magnetic grain trapping is so far the one and only mechanism proposed for prolonging the lifetime of hot grains in the vicinity of main-sequence stars, the refutal of the magnetic grain trapping leads us to the conclusion that the mechanism of trapping hot grains in the vicinity of a central star remains a mystery. The presence of hot dust grains has been identified among approximately 10--30\% of main-sequence stars by near-infrared interferometric observations \citep{absil-et-al2009,absil-et-al2013,ertel-et-al2014,nunez-et-al2017}. On the one hand, the detection of hot grains in the vicinity of main-sequence stars does not show clear dependences on stellar parameters nor spectral classes. On the other hand, a temporal variation in the presence of hot grains implies that an unsuccessful detection of hot grains in a specific observation does not guarantee the absence of hot grains in other periods of time \citep{ertel-et-al2016,nunez-et-al2017}. This does not conflict with a temporal variation in the detection of near-infrared excess around $4~R_\sun$ fron the G-type Sun, which has been attributed to the solar dust ring \citep[see][]{kimura-mann1998b}. These could be interpreted as the fact that the presence of hot grains in the vicinity of main-sequence stars is more ubiquitous among main-sequence stars than previously thought. | 18 | 8 | 1808.03389 |
1808 | 1808.01490_arXiv.txt | \cite{Zavarygin2018} compiled a list of 15 deuterium abundance measurements, discarded two because the remaining 13 measurements are then consistent with gaussianity, and found that the weighted mean baryon density ($\Omega_b h^2$) determined from the 13 measurements is mildly discrepant ($1.6\sigma$) with that determined from the Planck 2015 cosmic microwave background anisotropy data in a flat cosmogony. We find that a median statistic central estimate of $\Omega_b h^2$ from all 15 deuterium abundance measurements is a more accurate estimate, is very consistent with $\Omega_b h^2$ estimated from Planck 2015 data in a flat cosmogony, but is about $2\sigma$ lower than that found in a closed cosmogonical model from the Planck 2015 data. | \textbf{Introduction}} By the time the Universe was a few minutes old, the strong force had fused together neutrons and protons and synthesized the light nuclei. In the standard cosmological model the primordial light nuclei abundances depend only on $\Omega_b h^2$ (here $\Omega_b$ is the baryonic matter density parameter and $h$ is the Hubble constant in units of 100 km $\rm{s}^{-1}$ $\rm{Mpc}^{-1}$). Consequently, primordial light nuclei abundance measurements can be used to determine $\Omega_b$, and $\Omega_b$ determined using different light nuclei must agree, if the standard model is correct. Deuterium is particularly valuable in this regard as its predicted abundance is quite sensitive to the value of $\Omega_b h^2$. Spectroscopic analysis of absorption of quasar light by foreground low-metallicity gas clouds is used to estimate the primordial deuterium to hydrogen abundance, $\mathrm{(D/H)}_{\rm{p}}$.\footnote{The analysis focuses on Lyman absorption lines produced by the gas clouds. The deuterium Lyman lines are at slightly shorter wavelengths than those of hydrogen. The difference in absorption at these two sections of wavelengths determines D/H.} \cite{Zavarygin2018}, hereafter Z18, have compiled a list of 15 $\mathrm{(D/H)}_{\rm{p}}$ measurements. With the assumption of a cosmological model, $\Omega_b h^2$ can also be determined by fitting the cosmological model to observational data, such as cosmic microwave background (CMB) anisotropy measurements \citep{PlanckCollab2016b}. Comparisons between $\Omega_b h^2$'s determined from different data is a particularly compelling test of the standard cosmological model, and this can also be used to constrain other cosmological model parameters. Z18 argue that two of their 15 measurements are outliers, if their 15 measurements were drawn from a Gaussian distribution. Discarding these two measurements Z18 determine a weighted mean $\Omega_b h^2$ using the remaining 13 deuterium abundance measurements. They note that this value differs at $1.6\sigma$ from that determined by using the Planck 2015 TT + lowP + lensing CMB anistropy data \citep{PlanckCollab2016b}. Non-gaussian data compilations are not that rare \citep{Bailey2017}. Well-known examples include Hubble constant measurements \citep{Chen2003, Bethapudi2017, Zhang2018}, $^7 \rm{Li}$ abundance data \citep{Crandall2015a, Zhang2017}, LMC and SMC distance observations \citep{de Grijs2014, Crandall2015b}, and the Milky Way $R_0$ and $\Theta_0$ parameter measurements \citep{de Grijs2016, Camarillo2018a, de Grijs2017, Rajan2018, Camarillo2018b}. Since gaussianity is assumed in parameter estimation \citep[e.g.,][]{Samushia2007, Samushia2010, Farooq2015}, much effort has been devoted to testing for intrinsic non-gaussianity \citep[][and references therein]{Park2001, PlanckCollab2016a}, as distinct from non-gaussianity introduced by the measurement procedure. Conventional techniques cannot be used to analyze data with non-gaussian errors \citep{Gott2001, Bailey2017}; this was one of the motivations for the development of median statistics \citep{Gott2001}. Median statistics does not make use of the errors on individual measurements and so is not affected by incorrect errors. On the other hand, since it does not use this information it is less constraining than a weighted mean analysis. Perhaps the most well known example of the use of median statistics is its application to the measurement of the Hubble constant \citep{Gott2001, Chen2003, Chen2011, Rajan2018}. In this paper we apply median statistics to Z18's compilation of 15 $\mathrm{(D/H)}_{\rm{p}}$ measurements. We first examine the gaussianity of the Z18 data compilation. In agreement with Z18, we find that the full 15 measurements data set is non-gaussian, while their favored truncated set of 13 measurements is consistent with gaussianity. We then argue that the less precise median statistics summary estimate for $\Omega_b h^2$ for all 15 measurements is a more accurate representation of the data than is the more precise weighted mean summary estimate for the truncated data set of 13 measurements. We find that the median statistics $\Omega_b h^2$ determined from $\mathrm{(D/H)}_{\rm{p}}$ measurements is very consistent with those determined from other cosmological data in the context of spatially-flat cosmogonies, but is about 2$\sigma$ lower than those determined from cosmological data when using non-flat cosmogonies. | \textbf{CONCLUSION}} Our median statistics analysis of the complete set of 15 (D/H)$_{\rm{p}}$ measurements compiled by Z18 results in an $\Omega_b h^2$ estimate that is very consistent with those estimated from cosmological data in spatially-flat cosmogonies, but is about $2\sigma$ lower than what cosmology data favor in closed models. A full likelihood analysis including other cosmological data will need to be performed in order to determine the proper significance of this result. | 18 | 8 | 1808.01490 |
1808 | 1808.08052_arXiv.txt | Homotopy methods have been widely utilized to solve low-thrust orbital transfer problems, however, it is not guaranteed that the optimal solution can be obtained by the existing homotopy methods. In this paper, a new homotopy method is presented, by which the optimal solution can be found with probability one. Generalized sufficient conditions, which are derived from the parametrized Sard's theorem, are first developed. A new type of probability-one homotopy formulation, which is custom-designed for solving minimum-time low-thrust trajectory optimization problems and satisfies all these sufficient conditions, is then constructed. By tracking the continuous zero curve initiated by an initial problem with known solution, the optimal solution of the original problem is guaranteed to be solved with probability one. Numerical demonstrations in a three-dimensional time-optimal low-thrust orbital transfer problem with 43 revolutions is presented to illustrate the applications of the method. | \par The use of low-thrust propulsion in a variety of space missions~\citep{Rayman,Kawaguchi,Kugelberg} has gained great attention in the space community, which allows a substantial reduction of propellant consumption in virtue of its high specific impulse compared to traditional chemical propulsion. However, solving low-thrust trajectory optimization problems is known to be highly challenging, the solution methods of which are usually categorized as direct methods and indirect methods~\citep{Betts1998}. Direct methods convert the optimal control problems into nonlinear programming problems by appropriate discretization~\citep{Hargraves2,Armellin}, which are straightforward and robust to accommodate complex conditions. However, the optimality of the obtained solutions is not guaranteed. Indirect methods convert the original problems to two-point boundary-value problems~(TPBVPs) according to the optimal control theory, the solutions of which are guaranteed to be at least local extremals~\citep{Kechichian,Zeng2014Fast,Zhang2015Low,Jiang2016Systematic}. The main disadvantages associated with TPBVP are that its convergence domain is narrow and its solution is extremely sensitive to the initial unknowns if a single-shooting method is utilized, especially for the low-thrust trajectory optimization problems with long flight duration and many revolutions. Although multiple-shooting techniques exist that can efficiently enhance the robustness of the indirect methods, however, the number of the unknown variables may increase significantly~\citep{Betts1998,Taheri2017Co}. \par Homotopy methods, the principle of which is that a given problem is embedded into a family of problems parameterized by a homotopic parameter, and the optimal solution to the original problem is obtained by tracing the optimal solutions of the embedded problems~\citep{Watson2002}, have been widely applied to circumvent the above disadvantages of solving low-thrust trajectory optimization problems by single shooting indirect methods. References~\citep{Bertrand, Haberkorn2004Low, Gergaud2006Homotopy, Guo2012, Jiang2012, Zhang2015Low, Chen2016Optimality, Enhanced2016, Chi2017Homotopy,Zhao2017Target, Pan2018} have successfully utilized homotopy methods to solve minimum-fuel low-thrust orbital transfer problems, the optimal thrusts of which are discontinuous bang-bang controls. In these methods, the homotopic parameter is embedded into the performance index to provide continuous transition of optimal controls from the initial problem to the original one. It has been widely observed that the original fuel-optimal low-thrust trajectory optimization problems can be easily solved with probability one once the initial solutions are achieved~\citep{Guo2012, Jiang2012,Zhang2015Low, Chen2016Optimality, Chi2017Homotopy, Pan2018}. In contrast, homotopy methods for solving minimum-time low-thrust orbital transfer problems, whose optimal thrusts keep constant during the whole optimal trajectory, are still not satisfactorily developed. In Refs.~\citep{Caillau2003,Yue2010Indirect,Caillau2012Minimum}, the homotopy parameter is embedded into the thrust magnitude, and the minimum-time problem with sufficiently large thrust magnitude is taken as the initial problem for the homotopic approach. However, it is also not guaranteed that the optimal solution to the original problem can be obtained. \citep{Pan2016Double} presented a new double-homotopy method to construct {\it discontinuous} homotopy path which connects the initial and the original problem. However, the construction of discontinuous homotopy path is only valid under an assumption that multiple branches of homotopy path always exist at specific homotopic parameter, which may not be ensured for different occasions. Hence, the convergence to the optimal solution of the original problem by the double-homotopy method is still not guaranteed. Thus, homotopy methods, which construct {\it continuous} homotopy path to solve minimum-time low-thrust orbital transfer problems with probability one, are still unsettled. \par In this paper, a new probability-one homotopy method is presented to solve minimum-time low-thrust orbital transfer problems. Parametrized Sard's theorem~\citep{Chow1978Finding} and Watson's sufficient conditions~\citep{Watson2002} are first revisited, which ensure the probability-one convergence of the homotopy methods. Generalized sufficient conditions are then derived, and a new probability-one homotopy formulation is custom-designed to satisfy all the prerequisites of the generalized sufficient conditions for the minimum-time low-thrust orbital transfer problems. Numerical solutions of a minimum-time low-thrust orbital transfer problem are provided to demonstrate the effectiveness of the proposed method, the initial ratio of thrust-to-weight of which is as small as $6.8 \times {10^{ - 5}}$. | \label{Section:CONCLUSION} \par This paper presents a new probability-one homotopy method specifically for solving minimum-time low-thrust orbital transfer problems. The generalized sufficient conditions derived in this paper, reasonably explains the failure of the existing homotopy methods, in which the 4th prerequisite of the sufficient conditions is not satisfied, and also ensures the success of the proposed method. The optimal solution of the original problem can be easily found with probability one by tracing the continuous zero curve constructed by the proposed homotopy method. A new discovery is that the continuous zero curve exists even when the initial and the original problem have different number of orbital revolutions, which is not previously known in the literatures. Besides, in the literatures, it was concluded that the homotopy methods are only valid when the initial and the original problem share the same revolution number, which is also incorrect as pointed out in this paper. As illustrated in the numerical demonstrations, the optimal solution of the original problem with 43 revolutions, can be solved by starting from an initial problem with 1 revolution by the proposed method. Thus, the proposed homotopy method provides an efficient approach to find the optimal solution for the minimum-time low-thrust trajectory optimization problems, the convergence of which is probability one. \vspace{0.1in} \begin{center} {\bf Acknowledgments } \end{center} \par The authors gratefully acknowledge the support to this work by the National Natural Science Foundation of China (Grant No. 11672234). \vspace{0.2in} \clearpage | 18 | 8 | 1808.08052 |
1808 | 1808.04083_arXiv.txt | The cosmic microwave background (CMB) fluctuations effectively measure the basic properties of the universe during the recombination epoch. CMB measurements fix the distance to the surface of last scatter, the sound horizon of the baryon-photon fluid and the fraction of the energy density in relativistic species. We show that the microwave background observations can also very effectively constrain the thickness of the last scattering surface, which is directly related to the ratio of the small-scale E-mode polarization signal to the small-scale temperature signal. The current cosmological data enables a 0.1\% measurement of the thickness of the surface of last scatter: $19 \pm 0.065$ Mpc. This constraint is relatively model-independent, so it can provide a new metric for systematic errors and an independent test of the $\Lambda {\rm CDM}$ model. On the other hand, it is sensitive to models which affect the reionization history of the universe such as models with annihilating dark matter and varying fundamental constants (e.g., the fine-structure constant, $\alpha_{\rm EM}$, and electron rest mass, $m_{\rm e}$) and as such can be used as a viable tool to constrain them. | The cosmic microwave background (CMB) stores a tremendous amount of information about the early history of the Universe and its subsequent evolution. Its large-scale anisotropies indicate the existence of tiny fluctuations in the primordial gravitational potential that are the seeds for the formation of galaxies and other large-scale structure. Shortly after the discovery of the CMB in 1965, it was also shown \cite{1968ApJ...153L...1R} that anisotropic Thomson scattering of photons and electrons induces a degree of linear polarization in the data. The polarization properties of the CMB provide yet another set of observables, the measurement of which enriches greatly our understanding of the Universe. The CMB anisotropies were formed primarily around the epoch of hydrogen recombination around redshift $z \sim 1100$ defined by the peak of the visibility function, $v = |\dot \tau| e^{-\tau}$, which describes the probability that a CMB photon last scattered off free electrons at a particular point in the history of the Universe. The shape of the CMB power spectrum is thus most sensitive to changes around the peak of the visibility function. For example, the location of that peak determines the distance to the last scattering surface, which in turn determines the positions of the peaks of the CMB power spectra. If we were to increase the width of the visibility function, that would correspond to a prolonged period of recombination, leading to more Thomson scatterings of photons off free electrons. On scales smaller than the recombination width, these scatterings lead to the cancellation of the CMB anisotropies along the line of sight, while on larger scales they lead to enhancement of the polarization signal. Similarly, changes in the ionization history, the photon-baryon sound speed, the gravitational potential around matter-radiation equality and other primordial properties which we have no direct way of probing would likely lead to measurable changes in the CMB power spectra and the baryon acoustic oscillation (BAO) peak, which would then allow us to put constraints on the features of the early Universe \cite{harari}. For example, models which involve annihilation of dark matter to Standard Model particles between the period of recombination and reionization result in a modification of the ionization history, as they lead to a heating of the baryons and ionization of the neutral hydrogen, and thus alterations in the CMB visibility function. Another such example are models which predict variations of the fundamental constants -- e.g., the fine-structure constant, $\alpha_{\rm EM}$, and the electron rest mass, $m_{\rm e}$ provide another such example and can thus directly impact CMB observables \cite{1999PhRvD..60b3516K,2000PhRvD..62l3508A,2001PhRvD..63d3505B, 2001PhRvD..64f3505A,2004MNRAS.352...20R,2009arXiv0906.0329S, 2010HiA....15..307S,2012PhRvD..85j7301M}. Due to their nature of affecting mostly the large-scale observables, constraints on such models are not expected to become much more stringent with the drastic improvement in sensitivity expected of future cosmological surveys. Over the past decade, cosmologists have been exploring many alternative ways to test our understanding of the early Universe. A validated approach they have taken is to introduce new parameters into the analysis of cosmological data, which can serve as powerful probes for discrepancies with our predictions and can help in the detection of systematic errors in our measurements or problems with the $\Lambda {\rm CDM}$ model. In principle, to show that a given new parameter may be of such use, one needs to check that it is not correlated with the standard parameters, i.e. that it reflects a different physical effect. The degree of correlation between the new and the standard parameters can be tested by studying their contour plots in a Monte-Carlo engine analysis \cite{2013ascl.soft07002A}. An example of a parameter which probes the properties of the early Universe is the width of the visibility function during the ``last scattering'' of photons, By changing its width the strip of time from which the photons could have come would be broadened or narrowed. A longer period of last scattering would then result in a more polarized signal on large scales, since the photons would have scattered off the electrons a larger number of times \cite{harari,hubyharari}. Therefore, an intriguing question to explore is: how well can we constrain the width of the last scattering surface with the most recent polarization data from the $Planck$ team. So far, the width of the last scattering surface has not been measured directly by CMB experiments, but its theoretical value is readily computed to be $\Delta \eta_\ast \approx 19 \ \rm{Mpc}$ ($\Delta z_\ast \approx 90$) using the latest cosmological codes, e.g. CLASS \cite{2011JCAP...07..034B}, and standard values for the cosmological parameters from \textit{Planck} \cite{Ade:2015xua}. This paper is organized as follows. We first provide motivation for our work by stating a relationship between the polarization-to-temperature ratio of the power spectra and the width of the visibility function pointed out by \cite{harari}. We then parametrize this width through $\alpha_{\rm vis}$ and explore how varying this new parameter affects the observable power spectra. Finally, we put constraints on its value using the latest CMB data from $Planck$, study its degeneracy with other parameters, and discuss its potential as a model-independent test of the $\Lambda \rm{CDM}$ model and of systematic errors. It also provides a powerful tool to constrain models which alter the reionization history of the Universe such as dark matter annihilation and decay and variable $\alpha_{\rm EM}$ and $m_{\rm e}$ models. | \label{chap:conc} In this paper, we explored the effect of the width of the last scattering surface on the power spectrum and the constraints we can obtain on its thickness given our current data. We found that with our current measurements of the temperature and the large-scale polarization power spectra, the thickness of the last scattering surface can be constrained to an astounding precision: $\sigma [\Delta \eta_\ast] \approx \textbf{0.11 {\rm Mpc}}$. The constraint comes from the Silk damping tail, which gets suppressed as we broaden the width of the visibility function, and from the large-scale polarization, which is strongly dependent on the width. If we include the polarization data from $Planck$ 2015 on all scales, we get the much tighter constraint of $\sigma [\Delta \eta_\ast] \approx \textbf{0.065 {\rm Mpc}}$. This is a consequence of the fact that the polarization-to-temperature ratio is proportional to the width of the last scattering surface squared. We believe that constraining this parameter is a good way to test our current model and probe for physics beyond the Standard Model. In the near future, the new polarization data from upcoming experiments such as Simons Observatory (SO) and CMB-S4 \cite{Abazajian:2016yjj} should allow us to measure the parameter $\alpha_{\rm vis}$ with even greater precision. High-$\ell$ data from SPT and ACT should in principle also help measure better the diffusion damping tail and thus put constraints on $\alpha_{\rm vis}$. Including $\alpha_{\rm vis}$ when analyzing the new datasets will enable us to detect deviations from the $\Lambda {\rm CDM}$ model and look for systematic errors in these new datasets. In addition, it will be useful when constraining models which alter the reionization history of the Universe such as self-interacting dark matter models and variable-$\alpha$ models. | 18 | 8 | 1808.04083 |
1808 | 1808.06615_arXiv.txt | LSST will supply up to $10^6$ supernovae (SNe) to constrain dark energy through the distance--redshift ($D_L$--$z$) test. Obtaining spectroscopic SN redshifts (spec-$z$s) is unfeasible; alternatives are suboptimal and may be biased. We propose a powerful multi-tracer generalization of the Alcock-Paczynski test that pairs redshift-free distance tracers and an overlapping galaxy redshift survey. Cross-correlating $5\times 10^4$ redshift-free SNe with DESI or Euclid outperforms the classical $D_L$--$z$ test with spec-$z$s for all SN. Our method also applies to gravitational wave sirens or any redshift-free distance tracer. | 18 | 8 | 1808.06615 |
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1808 | 1808.03024_arXiv.txt | We have used the Australia Telescope Compact Array (ATCA) to conduct further observations of the 36.2-GHz ($4_{-1}\rightarrow3_0$E) methanol transition towards the nearby active galaxy NGC~4945. These observations have led to a more accurate determination of the offset between the maser emission and the nucleus of NGC~4945 with a typical synthesised beam of $6^{\prime\prime} \times 4^{\prime\prime}$ ($108\times72$ pc). This corresponds to a factor of 4 improvement with respect to the major-axis of the beam. Other transitions of methanol and lines of other molecular species were obtained alongside the 36.2-GHz methanol emission, with strong detections of HC$_3$N~(J~=~$4 \rightarrow 3$) and CS~(J~=~$1 \rightarrow0$) presented here. We do not detect thermal methanol (5$\sigma$ upper limit of 5~mJy in a 6~\kms\ channel) from the 48.4-GHz ($1_{0}\rightarrow0_0$A$^+$) ground-state transition, nor emission from the 44.1-GHz ($7_{0} \rightarrow 6_1 $A$^+$) class~I maser transition (5$\sigma$ upper limit of 6 mJy in a 3~\kms\ channel). We also present a comparison of the class~I maser emission observed towards NGC~4945 with that from NGC~253 and towards the Galactic giant molecular cloud G~1.6-0.025. | \begin{table*} \begin{center} \caption{Details of the observed transitions in NGC~4945. All rest frequencies were taken from the online database: \textit{NIST Recommended Rest Frequencies for Observed Interstellar Molecular Microwave Transitions by Frank J. Lovas}\protect\footnote[1]{Test 1, 2 3}.} \begin{tabular}{llcllcl} \hline \multicolumn{1}{c}{\bf Species} & \multicolumn{1}{c}{\bf Transition} & \multicolumn{1}{c}{\bf Rest} & \multicolumn{1}{c}{\bf Array} & \multicolumn{1}{c}{\bf Epoch} & \multicolumn{1}{c}{\bf Integration} &\multicolumn{1}{c}{\bf Detection} \\ & & \multicolumn{1}{c}{\bf Frequency} & \multicolumn{1}{c}{\bf Configuration}& & \multicolumn{1}{c}{\bf Time} & \\ & & \multicolumn{1}{c}{(GHz)} & & & \multicolumn{1}{c}{(min)} & \\ \hline Methanol & $4_{-1} \rightarrow 3_0 $E & 36.169265 & EW352, H168, H214 & 2015 Aug, 2017 Jun, 2017 Oct & 265.2 & detection \\ & $7_{-2} \rightarrow 8_{-1} $E & 37.703696 & EW352, H168, H214 & 2015 Aug, 2017 Jun, 2017 Oct & 265.2 & non-detection \\ & $7_{0} \rightarrow 6_1 $A$^+$ & 44.069476 & H214 & 2017 Oct & 178.2 & non-detection \\ & $1_{0} \rightarrow 0_0 $A$^+$ & 48.372456 & H214 & 2017 Oct & 138.2 & non-detection \\ HC$_3$N & $J=4 \rightarrow 3~F = 4 \rightarrow 3$ & 36.392332 & H168, H214 & 2017 Jun, 2017 Oct & 197.4 & detection \\ HC$_5$N & $J=12 \rightarrow 11$ & 31.951777 & H168, H214 & 2017 Jun, 2017 Oct & 197.4 & non-detection \\ NH$_3$ & $12_{12} \rightarrow 12_{12}$ & 31.424943 & H168, H214 & 2017 Jun, 2017 Oct & 197.4 & non-detection \\ & $13_{13} \rightarrow 13_{13}$ & 33.156849 & H168, H214 & 2017 Jun, 2017 Oct & 197.4 & non-detection \\ CH$_3$CN & $2_{0} \rightarrow 1_{0},~F=3\rightarrow2$ & 36.795568 & H168, H214 & 2017 Jun, 2017 Oct & 98.4 & non-detection \\ SiO & $J=1 \rightarrow 0,~\text{v}=0$ & 43.423853 & H214 & 2017 Oct & 178.2 & detection \\ CS & $J=1\rightarrow0$ & 48.990955 & H214 & 2017 Oct & 138.2 & detection \\ \hline \end{tabular} \label{tab:observations} \end{center} \end{table*} Methanol maser emission is divided into two classes based on pumping mechanism. Methanol masers pumped via collisional processes are defined as class~I, while those that are radiatively pumped are considered class~II \citep{Batrla+87, Menten91a}. Both classes of methanol masers are commonly observed throughout the Milky Way, with over 1200 unique sources discovered \citep[e.g][]{Ellingsen+05,Caswell+10,Caswell+11, Voronkov+14, Breen+15, Green+10, Green+12a, Green+17}. In contrast with this, outside of our Galaxy we have relatively few detections of methanol maser emission. Extragalactic class~II masers have been detected in the Large Magellanic Cloud and M31 \citep{Green+08,Ellingsen+10,Sjouwerman+10}. These extragalactic class~II masers appear to be extremely luminous examples of their Galactic counterparts. Conversely, the extragalactic class~I methanol masers are not yet a well understood phenomenon, with observed emission unable to simply be considered large-scale emission from Galactic-style class~I masers. Currently there are 6 reported examples of class~I maser emission towards extragalactic sources, 36.2-GHz emission in NGC~253, Arp~220, IC~342, NGC~6946 and NGC~4945 \citep{Ellingsen+14, Chen+15, McCarthy+17, Gorski+18} and 84.5-GHz emission in NGC~1068 \citep{Wang+14}. Of these 6 sources, only NGC~253 has been detected in multiple epochs with multiple telescopes \citep{Ellingsen+14,Ellingsen+17b,Chen+18,Gorski+18}. % Class~I methanol emission is a powerful tool for understanding star-formation within our Galaxy. More than 600 unique sources of class~I maser emission are observed within the Milky Way \citep[e.g.,][]{Slysh+94,Valtts+00,Ellingsen+05,Chen+11,Gan+13,Jordan+17}. Galactic class~I masers are generally associated with shocked gas driven by the expansion of \ionhy regions or molecular outflows \citep{Kurtz+04, Cyganowski+09, Cyganowski+12, Voronkov+10a, Voronkov+14}. However, it is not yet known if, or how, these highly luminous extragalactic class~I masers relate to the star-formation of their host galaxies. In Galactic star formation regions the two most commonly observed class~I methanol maser transitions are the 36.2- and the 44.1-GHz, with the latter generally being the stronger of the two \citep[e.g. ][]{Voronkov+14}. \citep{Ellingsen+17b} detected weak 44.1-GHz methanol maser emission associated with two of the 36.2-GHz sites in NGC\,253. They suggest that the low intensity of the 44.1-GHz transition compared to the 36.2-GHz in NGC\,253 is strong evidence that extragalactic class~I methanol maser emission cannot be explained as being due a large number of Galactic-like star formation regions within a small volume, but is rather a new and different extragalactic masing phenomenon. It appears that the extragalactic variants may evolve from large-scale molecular inflow inside their host galaxies \citep{Ellingsen+17b}. However, this has so far only been verified in the case of NGC~253. Developing an understanding of the pumping environments responsible for this phenomenon is one of the most important factors in determining its usefulness as a probe of galactic properties. NGC~4945 is a nearby \citep[assumed distance of 3.7$\pm0.3$\,Mpc;][]{Tully+13} spiral galaxy, with a hybrid AGN and starburst nucleus. The starburst is the primary source of energy for exciting photo-ionized gas, due to heavy obscuration of the AGN by dust \citep{Marconi+00, Spoon+00, Spoon+03, Perez-Beauputis+11}. The star-formation rate in NGC~4945 is more than three times that of the Milky Way \citep[$4.35\pm0.25$ M$_{\sun}$ yr$^{−1}$ for the nuclear region of NGC~4945 only, compared to $1.65\pm0.19$ M$_{\sun}$ yr$^{−1}$ for the entire Milky Way;][]{Bendo+16,Licquia+15} and approximately 20 percent higher than that of the similar (in terms of galactic properties and maser luminosity) extragalactic class~I maser host galaxy, NGC~253 \citep{Strickland+04}. NGC~4945 is also host to various transitions of water megamasers, predominantly located in a circumnuclear accretion disk \citep{Greenhill+97, Hagiwara+16,Humphreys+16, Pesce+16}. We have undertaken new observations of the 36.2-GHz methanol maser transition in NGC\,4945 to better determine its location with respect to the host galaxy and other molecular gas. In addition to the 36.2-GHz methanol transition we have also observed the 44.1-GHz class~I methanol maser transition to determine if the NGC\,4945 shows a similar pattern to NGC\,253 with this transition being relatively much weaker than is observed towards Galactic class~I methanol masers associated with high-mass star formation regions. We were able to include observations of a number of thermal molecular transitions simultaneously with the maser observations and we present the results of those observations and compare them with the recent, sensitive high-resolution molecular line ALMA observations at 3-mm made by \citet{Henkel+18}. We currently have a sample of only six known extragalactic class~I methanol maser sources and by obtaining a range of complementary spectral line and other data and comparing the results for NGC\,4945 with other sources we hope to improve understanding of this new phenomenon and its relationship to the properties of the host galaxy. \footnotetext[1]{https://physics.nist.gov/cgi-bin/micro/table5/start.pl} | Our high-resolution imaging follow up of the 36.2-GHz class~I methanol maser emission in NGC~4945 has confirmed its offset position from the galactic nucleus. Assuming the region is part of the disk, it would be located at a galactocentric radius of approximately 670~pc and is likely associated with the interface region between galactic bar and south-eastern spiral arm on the front side of the galaxy. This position corresponds to the position of a hotspot observed in dense gas tracers, likely indicating an association between the masing region and a giant molecular cloud. We detect methanol emission from neither the 44.1-GHz masing transion nor the 48.4-GHz thermal transition towards NGC~4945 and this indicates a high optical depth for the 36.2-GHz class~I masing region. The 7-mm continuum source is offset by $0\farcs8\pm0\farcs4$ to the north-west of the dynamical centre of NGC~4945 ($\alpha_{2000} = 13^{\text{h}}~05^{\text{m}} ~ 27^{\text{s}}.467\pm0^{\text{s}}.032$ and $\delta_{2000} = -49^\circ ~ 28^\prime ~ 04\farcs8\pm0\farcs3$). Emission from the HC$_3$N $J = 4 \rightarrow 3$, CS $J = 1 \rightarrow 0$ and SiO $J = 1 \rightarrow 0$ transitions were also detected towards NGC~4945. All of these thermal transitions were observed towards the central region of NGC~4945 (consistent with the 3-mm ALMA observations). Additionally, all of these species display strong absorption towards the continuum source, with a peak absorption component at 636~\kms. None of these molecular species were detected towards the offset location where the methanol masers are observed. We identify many similarities between the class~I methanol masers of NGC~4945 and NGC~253. There is the possibility that the optical depth of the maser in NGC 4945 is higher, leading to particularly high intensity ratios between the 36.2- and 44.1-GHz methanol lines. A comparison with the giant molecular cloud G\,1.6-0.025 revealed similar line ratios between the 36.2- and 44.1-GHz methanol maser transitions, though unlike G\,1.6-0.025, the maser region in NGC~4945 may be associated with a region of enhanced star-formation. | 18 | 8 | 1808.03024 |
1808 | 1808.06092_arXiv.txt | New {\em Chandra}\/ High Resolution Camera pointings on the ``non-coronal'' red giant Arcturus (HD\,124897; $\alpha$~Boo: K1.5~III) corroborate a tentative soft X-ray detection in a shorter exploratory exposure sixteen years earlier. The apparent source followed the (large) proper motion of the nearby bright star over the intervening years, and there were null detections at the previous location in the current epoch, as well as at the future location in the earlier epoch, reducing the possibility of chance coincidences with unrelated high-energy objects. The apparent X-ray brightness at Earth, averaged over the 98~ks of total exposure and accounting for absorption in the red giant's wind, is $\sim 2{\times}10^{-15}$ erg cm$^{-2}$ s$^{-1}$ (0.2--2~keV). Systematic errors in the energy conversion factor, devolving from the unknown spectrum, amount to only about 10\%, smaller than the ${\sim}$30\% statistical uncertainties in the count rates. The X-ray luminosity is only $3{\times}10^{25}$ erg s$^{-1}$, confirming Arcturus as one of {\em Chandra's}\/ darkest bright stars. | Arcturus ($\alpha$~Bo\"otis; HD\,124897; K1.5~III)\footnote{Unless otherwise stated, stellar properties are from SIMBAD.} is the brightest star at northern declinations, third brightest overall. It is an old, solar mass, slightly metal poor red giant, only 11~pc from the Sun (e.g., Ram{\'{\i}}rez \& Allende Prieto 2011). Arcturus is of interest, among other reasons, because it mirrors the evolutionary fate that awaits the Sun some 5 billion years hence. Another curiosity is that the red giant belongs to a well-defined stellar stream (the eponymous Moving Group): possibly the remnants of an ancient dissolved open cluster; more speculatively a tidal tail stripped from a satellite galaxy that wandered too close to the Milky Way eons ago (e.g., Navarro et al.\ 2004); or perhaps simply a dynamical resonance in the Galactic disk (e.g., Bensby et al.\ 2014). In the early days of X-ray astronomy, low-mass ($\gtrsim 1\,M_{\odot}$) red giants like Arcturus rarely were detected in high-energy surveys by pioneering observatories like {\em Einstein}\/ (Vaiana et al.\ 1981; Ayres et al.\ 1981), and later {\em R\"ontgensatellit} ({\em ROSAT}\,) (Haisch et al.\ 1991). In contrast, yellow giants in the Hertzsprung gap, such as Capella ($\alpha$~Aurigae: G1~III + G9~III), often were strong coronal ($10^6$--$10^7$~K) emitters. In fact, Linsky \& Haisch (1979) earlier had proposed -- based on ultraviolet proxies -- a dividing line in the giant branch, separating the coronal ``haves" from the ``have nots." The ``non-coronal'' side (redward of spectral type K1~III) later became known as the ``coronal graveyard,'' after a deep X-ray pointing by {\em ROSAT}\/ failed to detect Arcturus, prototype of the class (Ayres, Fleming, \& Schmitt 1991). It was well known that an internal magnetic ``Dynamo'' -- relying heavily on stellar rotation -- underpins the cycling activity of sunlike stars, so the demise of coronae among the bloated, slowly spinning red giants seemed sensible. A decade later, in mid-2002, one of the new-generation X-ray facilities, {\em Chandra,}\/ turned its sharper gaze on Arcturus (Ayres, Brown, \& Harper 2003 [ABH]). In a 19~ks exposure with the High Resolution Camera (HRC-I), a mere 3 counts were recorded in a small detect cell centered at the predicted coordinates of the red giant. Nevertheless, thanks to unusually low cosmic background conditions, the few counts represented a moderately significant detection. The estimated X-ray luminosity of Arcturus was more than an order of magnitude lower than that of the average Sun, itself a rather mild coronal source. The $L_{\rm X}$ was especially diminutive given that the surface area of the K giant is more than 600 times that of the G dwarf. At about that time, the high-sensitivity ultraviolet spectrographs of {\em Hubble Space Telescope}\/ uncovered unexpected clues to the apparent coronal disappearing act. The first surprise was the clear presence of coronal proxy \ion{C}{4} 1548~\AA\ in non-coronal giants like Arcturus (Ayres et al.\ 1997), albeit weak enough to have escaped previous notice. \ion{C}{4} forms at $10^5$~K, hot enough that magnetic heating must be involved. The second surprise was that other hot lines, \ion{Si}{4} 1393~\AA\ and \ion{N}{5} 1238~\AA, showed stationary, sharp absorptions from cool species such as \ion{Ni}{2} and \ion{C}{1} (ABH). This implied that the hot emitting structures must be buried under a large overburden of lower temperature ($\sim 6000$~K) chromospheric material, a ``cool absorber'' if you will. The large column can suppress soft X-rays, but still pass FUV radiation. If the ``buried corona'' conjecture is correct, deep-seated magnetic activity on the non-coronal giants might be responsible for stirring up their atmospheres and initiating their powerful winds: $10^4$ times the Sun's mass loss rate (for the specific case of Arcturus), but much cooler than the solar coronal counterpart, $T\sim 10^4$~K versus $\sim 10^6$~K (e.g., O'Gorman et al.\ 2013). The red giant outflows are important to galactic ecology, but the motive force behind the winds has remained elusive. The best test of the buried corona hypothesis would be an X-ray spectrum, to judge the extent of the putative chromospheric soft absorption. A minimal CCD-resolution ($E/\Delta{E}\sim 50$) energy distribution in the 0.25--10~keV band generally would require $\sim 10^3$ net counts; out of the question with contemporary instrumentation, at least given the apparent faintness of the Arcturus source in 2002. However, recently the {\em Chandra}\/ Observatory offered a special opportunity to carry out observations that might help inform the design of next-generation X-ray facilities. A proposal for a deeper HRC-I exposure of Arcturus -- in essence a feasibility assessment for a future spectrum -- was among the projects chosen. Here, the results of the new X-ray observations of Arcturus are described, and their implications discussed. To preview the more detailed conclusions presented later, a source at the predicted location of Arcturus (accounting for proper motion) was clearly present in each of the two new observations; and especially the sum (including also the earlier [2002] pointing). | {\em Chandra}\/ pointings on the archetype non-coronal red giant Arcturus have secured moderately significant to very significant detections of an X-ray source at the stellar coordinates, in three epochs. Accidental X-ray objects at the two distinct locations in 2002 and 2018 (well-separated thanks to high proper motion of the bright star) are unlikely, given the lack of significant sources at the future and past positions in the respective epochs. Further, the high spatial resolution of {\em Chandra}\/ naturally minimizes source confusion. Although the apparent Arcturus X-ray source is rather faint, it nevertheless suggests that a future high-energy observatory with $\sim$100 times the contemporary {\em Chandra}\/ sensitivity, and similar or better spatial resolution, could collect a diagnostically valuable spectrum in a reasonable exposure ($\sim$100~ks). Such observations could help assess the properties of possibly buried coronae in the extended outer atmospheres of red giants, and perhaps also contribute to resolving the puzzle of their enigmatic winds. | 18 | 8 | 1808.06092 |
1808 | 1808.01167_arXiv.txt | In this study we derive the magnetized Taub adiabat equations from the hydrodynamic conservation conditions. We employ the magnetized taub adiabat equations to study the evolution of a magnetized neutron star to magnetized quark star. The pressure of the burnt quark matter has a maximum which indicates a bound on the maximum mass of the quark star. The central density of the neutron star and the angle between the rotation axis and the magnetic axis (defined as the tilt angle) are seen to be significant in determining the magnetic field, and the tilt of the quark star. The magnetic field and the tilt of the quark star can have a observational significance and can help in understanding the physics at high density and strong magnetic field. | Shock fronts are generally depicted by a discontinuous change in the characteristics of the medium which propagates faster than the speed of sound in that medium. In most plasmas, the width of the shock front is very thin, and it is usually considered to be a one-dimensional plane of discontinuity \cite{landau}. Taub \cite{taub} was the first to study the relativistic hydrodynamic shocks. He used the mass, momentum, and energy conservation laws to derive the relativistic Rankine-Hugoniot (RH) conditions. De-hoffmann \& teller \cite{hoffmann} displayed the theoretical treatment of hydrodynamic shocks in the presence of a magnetic field, which was followed by an avalanche of theoretical studies \cite{landau,cabannes}. However, the relativistic treatment of magnetized hydrodynamic shocks was first done by Lichnerowicz \cite{lichnerowicz,webb}. Other important works in this field were successively carried out by Appl \& Camezind \cite{appl}, Majorana \& Anile \cite{majorana} and Ballard and Heavens \cite{ballard} to name a few. The interaction between hydrodynamic motion and magnetic fields in conducting plasmas are essential in the problem of astrophysics, high-energy collision, and geophysics. Two individual cases of magneto-hydrodynamic waves are common in physics; the hydrodynamic shock and the electromagnetic wave. As the electromagnetic waves travel at the speed of light, we need to treat the problem relativistically. De-hoffmann \& Teller \cite{hoffmann} did that and treated the conducting fluid to be of having infinite conductivity. It was done by transforming the shock to a frame where the flow velocity is parallel to the magnetic field. This assumption prevents the self-induction of the magnetic field if the fluid is at rest, and is well suited for astrophysical scenarios as the spatial dimension of most of the astrophysical problems are very large. The standard technique of writing the jump conditions is to set the divergence of the stress-energy tensor to be zero and use the Gauss's theorem to get the jump conditions across the shock front. The conditions give the general mass, momentum, and energy continuity equations across the front. The three non-magnetized hydrodynamic equations can be expressed as a single equation known as Taub adiabat (TA) equation \cite{taub,thorne}. The equation connects the thermodynamic variables of one side of the front with variables on the other side, and is deprived of any velocity term. The astrophysical problem which frequently deals with hydrodynamic equations is the PT from a neutron star (NS) to a Quark star (QS). It was conjectured that at the center of NSs where supra-nuclear densities are believed to be present, normal hadronic matter (HM) is not the stable state of matter \cite{itoh,bodmer,witten}. At such densities HM is vulnerable to QM (quark being their constituent particles). Therefore, it is likely that a PT occurs taking the strongly interacting confined phase (HM) to a deconfined phase (QM). If the QM prevails over a substantial region at the core of a NS it is referred as a hybrid star and if the entire star consists of deconfined quarks it is known as a strange star. Such PT or conversion in NS is likely to liberate a significant amount of gravitational energy, which can power gamma-ray bursts \cite{berezhiani,mallick-sahu} and can have gravitational wave signals \cite{mallick-mag1,mallick-apj}. The initiation of this PT can be due to a variety of reasons: starting from cosmological quark nugget \cite{alcock}, mass-accretion \cite{alcock} to pulsar glitches \cite{chubarian}. The process of phase transition is also widely debated in literature: whether it is a detonation or a deflagration. One of the most usual ways of studying the PT is by employing the hydrodynamic equations. Cho et al. \cite{cho} were probably the first ones to use the hydrodynamic jump conditions to argue that weak detonation is a possible mode of combustion. However, Tokareva et al. \cite{tokareva} and Lugones et al. \cite{lugones} argued that the possible mode of combustion could even be fast detonation. Bhattacharyya et al. \cite{bhat,mallick-igor} argued that it can be a detonation or a deflagration depending on the density, whereas Drago et al. \cite{drago} demonstrated that it is always a deflagration if the process is exothermic. A detailed discussion on such scenario can be found in a recent paper by Furusawa et al \cite{furusawa}. Most of the studies discussed above employ the relativistic and non-relativistic RH equations; however, TA equation is not much studied. The magnetized counterpart of the TA is still to be derived and analyzed. In this article, we calculate the magnetized version of the TA equation and employ them to study the astrophysical scenario of PT. The usefulness of this equation lies in the fact they do not involve matter velocities and only deals with the thermodynamic variables of the state like pressure, density and energy. The paper is arranged as followed: In section 2 we calculate the magnetized TA equations from the conservation equations. Next, in section 3 we discuss the EoS and the magnetic field configuration of the star. Section 4 is dedicated to our results, and finally, in section 5 we summarize our work and draw conclusion from them. | In this article, we have derived the MTA equations, which was never realized before. These equations are different from the general conservation equations since the velocity terms are absent. The matter velocities can be calculated from the thermodynamic variables and magnetic fields. QM TA shows the same retracing nature even in magnetized plasmas. The pressure curve obtained from solving the MTA equations has a peak, which is reflected from the retracing of the TA of QM. The retracing nature of the TA and the maximum value of pressure of the QM indicates a mass bound to the QSs assumed to be formed by first order PT. \begin{figure} \includegraphics[width = 3.5in]{vel-mag.eps} \caption{(Color online) The upstream ($v_1$) and downstream ($v_2$) velocities are shown as a function of $\rho_b$ for input HM EoS. Both $v_1$ and $v_2$ first increases till a point then decreases and goes to zero. $v_1$ is always greater than $v_2$ however they reach the maximum point and goes to zero at same $\rho_b$ values. The velocities for NL3 EoS goes to zero at much smaller value of $\rho_b$ than for PLZ velocities. The velocities attain high non zero values at much higher densities.} \label{fig-vel} \end{figure} The angle between the matter velocities and the shock front along with the density of the NS is critical in determining the magnetic field of the burnt star and the downstream shock velocity angle. This can have significant observational consequence as it could determine whether the PT to a QS would result in a star less or more magnetic than the initial NS. The initial tilt angle (angle between the rotational axis and the magnetic axis) is also the angle between the shock front and matter velocity, assuming that the shock spreads spherically in the star. The final tilt of the QS can be different from the initial inclination of the NS. For instance, if a NS of about $1.2-1.4$ solar mass with small tilt angle suffers a PT the magnetic field of the QS would be larger than the initial NS. However, if the tilt angle is large, the QS has similar field strength as that of the NS. The situation is different for a more massive star, where a PT with small tilt angle would result in a QS whose magnetic field is less than the NS. The burning mechanism of the star is dependent on the magnetic field, the tilt angle and the density of the NS. A star of about $1.2-1.5$ solar mass with a moderate tilt angle is likely to undergo a detonation ($v_1 > v_2$) however a massive star of about $1.8-2$ solar mass with small tilt angle is expected to experience a PT via a deflagration process. There are some stars which are not prone to PT as indicated by the zeros of the velocity curve. On the other hand, a low sized star with high tilt angle is likely to undergo a PT via deflagration mechanism. To summarize, the MTA equations can be a tool to study PT in magnetized NSs. The results can be generalized and are true for hM and QM EoS. We are in the process of further analyzing this mechanism and studying other general features of the MTA. | 18 | 8 | 1808.01167 |
1808 | 1808.00357_arXiv.txt | {Light but massive cosmological neutrinos do not cluster significantly on small scales, due to their high thermal velocities. With finite masses, cosmological neutrinos become part of the total matter field and contribute to its smoothing. Structure formation in the presence of massive neutrinos is therefore impeded compared to that in the standard $\Lambda$CDM cosmology with massless neutrinos. Neutrinos' masses also distort the anisotropy power spectrum of cosmic microwave background (CMB). Furthermore, a finite chemical potential $\mu$ for cosmological neutrinos, still allowed by current data, would have a non-negligible impact on CMB and structure formation. We consistently evaluate effects of neutrino masses and chemical potentials on the matter power spectrum by use of a neutrino-involved N-body simulation, with cosmological parameters obtained from a Markov-Chain Moonte-Carlo (MCMC) refitting of CMB data. Our results show that while a finite averaged neutrino mass $m_\nu$ tends to suppress the matter power spectrum in a range of wave numbers, the neutrino degeneracy parameters ${\xi_i \equiv \mu_i /T}$ ($i=$1, 2, 3) enhance the latter, leading to a large parameter degeneracy between $m_\nu$ and $\xi_i$. We provide an empirical formula for the effects on the matter power spectrum in a selected range of wave numbers induced by $m_\nu$ and $\eta \equiv \sqrt{\sum_i \xi^2_i}$. Observing a strong correlation between $m_\nu$ and $\eta$, we propose a single redshift-independent parameter $m_\nu - \frac{4}{3}\eta^2$ to characterize the neutrino effects on the matter power spectrum.} \begin{document} | \label{sec:intro} Cosmological neutrinos are believed to be the most abundant fermions in the Universe. However, this cosmological neutrino background (CNB) is difficult to detect directly, due to its low temperature of about 2 K now. On the other hand, results of neutrino oscillation experiments show that there are three neutrino mass eigenstates, $m_{\nu_i} ( i =$ 1, 2, 3), with mass-squared splittings $\Delta m^2_{12} \approx 7.37 \times 10^{-5} $ eV$^2$ and $|\Delta m^2_{23}| \approx 2.50 \times 10^{-3} $ eV$^2$ \cite{nu_osc}. These slightly massive neutrinos effectively behave as hot dark matter, which has been ruled out as the dominant part of the dark matter \cite{hdm}, although it is still the only known component. This CNB is expected to have impact on the CMB anisotropies and large-scale structure (LSS) formation. Compared to those in the standard $\Lambda$CDM cosmology with massless neutrinos, the radiation-matter equality time and the Hubble expansion rate would be modified with massive cosmological neutrinos. The photon diffusion scale $\theta_d$ and sound horizon scale $\theta_s$ would then be altered, resulting in a modified angular anisotropy power spectrum of the CMB. The integrated Sachs-Wolfe effect and the gravitational lensing of CMB polarization are also affected by massive neutrinos via early structure growth \cite{aba_nu}. Altogether the updated constraints on the sum of neutrino masses are $\sum_i m_{\nu_i} < 0.23$ eV from the Planck 2015 observation \cite{planck2015}, and $\sum_i m_{\nu_i} < 0.12$ eV from the newest Planck 2018 result \cite{planck2018}. The large scale structure formation is more sensitive than CMB to the sum of neutrino masses. It is well known that the growth of structure is governed by the competing effects of the cosmic expansion and self-gravity of matter perturbations, both of which are affected by the massive neutrinos \cite{ab2013}. Unlike photons, slightly massive cosmological neutrinos experience a transition from being relativistic to non-relativistic as the universe cools down, resulting in a different background expansion rate from the $\Lambda$CDM model. On the other hand, compared to cold dark matter (CDM), cosmological neutrinos have a higher thermal velocity and are less likely to be trapped by potential wells. Thus their spatial distribution is much more disperse than CDM. As a consequence, below its free-streaming scale, the CNB slows down the growth of structures, leading to a suppression of the total matter power spectrum. This free-streaming effect has traditionally been studied in the linear regime \cite{cpma, eh}, and included in numerical Boltzmann codes such as \texttt{CAMB} \cite{camb}. However, as the Lagrangian perturbation theory is limited to the condition that the over-density field $|\delta| \ll 1$, this linear method cannot be used to explore the influence of massive neutrinos on the late non-linear growth of structures. Instead, cosmological N-body simulation or simulation-based Halofit formula should be used \cite{halofit}. A natural way to implement neutrinos in N-body simulations is to treat them as an independent kind of simulation particles, with a much higher typical velocity compared to that of CDM particles \cite{brandbyge08, viel10, vnavarro12, tiannu}. This particle-based simulation is in principle accurate but computationally expensive, and the consequent power spectrum is dominated by shot noise in small scales due to the finite number of neutrino particles \cite{ab2013, brandbyge08}. Another much cheaper alternative is to implement neutrinos as an over-density field on regular grids, and its evolution is then studied by linear perturbation theory. This method is justifiable since neutrinos do not significantly cluster below its free-streaming scale, $k_{fs}$, which is larger than the non-linear scale $k_{nl}$ \cite{ab2013}. This grid-based simulation was first proposed in \cite{brandbyge09}, in which the neutrino power spectrum was evolved by \texttt{CAMB} and improved to include the non-linear effect of CDM clustering in \cite{ab2013}. The consistency between these two approaches to include neutrinos in cosmological N-body simulation has been well tested in both \cite{ab2013} and this paper. Most of the studies mentioned above use the same cosmological parameter set when comparing the matter power spectra for different neutrino masses. However, the cosmological parameters obtained from fitting the CMB power spectrum also depend on the neutrino mass given. A fully consistent study should take this into consideration. This is regarded as a third mechanism for massive neutrinos to affect the LSS in our work, although it is not really independent from either the expansion history or the free-streaming effect. Because both CMB and LSS are sensitive to the sum of neutrino masses instead of the mass hierarchy \cite{hierarchy}, we add a new variable $m_{\nu} \equiv \sum m_{\nu_i} / 3$, the averaged mass for the three neutrino mass eigenstates, into a Markov-Chain Monte-Carlo code \texttt{CosmoMC} for fittings of cosmological parameters from the Planck 2015 CMB data (as the likelihood code of Planck 2018 is not released to public yet) \cite{planck2015, cosmomc}. N-body simulations are then generated using these sets of refitted cosmological parameters. In this way the study of neutrinos' influence on LSS is self-consistent. Another mystery about neutrinos is whether they are Majorana or Dirac particles. If we assign the chemical potentials of neutrinos to be $\{\mu_i\}$, where $i$ labels the neutrino mass eigenstates, then for anti-neutrinos they are $\{-\mu_i\}$. Naturally Majorana neutrinos must have $\mu_i\ = 0$, and if $\mu_i\ \neq 0$, neutrinos are Dirac fermions. Because the neutrino distribution has been frozen after decoupling, $\{\xi_i \equiv \mu_i / T\}$ are fixed and denoted as the neutrino degeneracy parameters. The difference between $\{\xi_i\}$ and $\{-\xi_i\}$ leads to an asymmetry in the neutrino and anti-neutrino number densities. Big Bang nucleosynthesis (BBN) constrains this asymmetry of electron-type neutrinos to be small, with $|\xi_e| \leq \mathcal{O} (10^{-2})$, while the total neutrino asymmetry is mainly constrained by the effective number of relativistic species $N_{eff}$ by CMB, and $\xi_{\mu, \tau}$ of $\mathcal{O}(1)$ is still allowed \cite{bbn, barenboim_asym}. In this paper we also include the possibility of finite $\{\xi_i\}$ in our \texttt{CosmoMC} fitting for cosmological parameters as well as the modified N-body simulation to study its influence on LSS, in addition to the effect of neutrino mass. We follow \cite{barenboim_mass_eigenstate} to set $\xi_e = 0$ and $\xi_\mu = \xi_\tau$, the latter because of the strong mixing between $\nu_\mu$ and $\nu_\tau$. Only one free parameter is left for $\{\xi_i\}$, which we choose to be $\eta \equiv \sqrt{\sum \xi_i^2}$. In this work, we examine the influence of the averaged neutrino mass $m_\nu$ and the neutrino degeneracy parameter $\eta$ on the total matter power spectrum. All three mechanisms: the modification of the cosmic expansion rate, the neutrino free-streaming effect and the shifts in cosmological parameters obtained from CMB fitting are consistently included. We mainly use the grid-based method of including neutrinos in N-body simulation, by our own modified version of the \texttt{Gadget2} code \cite{gadget2}. This paper is organized as follows. In Section 2 we elaborate on the calculation of neutrino energy density and the treatment of $\{\xi_i\}$ with known constraints. The introduction to our \texttt{CosmoMC} refitting is also included. We elaborate on the detailed procedure to conduct both the particle-based and grid-based simulations, and compare the measured total matter power spectra to show the consistency of these two methods in Section 3. We present our results in Section 4, and provide an empirical formula for the neutrino induced change in the matter power spectrum. Our summary and disscussion are in Section 5. | In this paper we have studied the impact of massive and degenerate cosmological neutrinos on the total matter power spectrum. Apart from the neutrino free-streaming effect and modified expansion history, we have also used the \texttt{CosmoMC} code to refit the cosmological parameters from the Planck 2015 CMB data when finite neutrino mass and degeneracy are allowed, so that the study is self-consistent. Our simulations show that for a reasonable range of parameter values, the neutrino degeneracy parameter squared $\eta^2$ has a comparable effect on the matter power spectrum as that of the mass parameter $m_{\nu}$, but with opposite signs, and so there could be an issue of parameter degeneracy. Thus previous studies that estimate the neutrino mass purely from the free-streaming effect on LSS may not be accurate and should be re-evaluated. We provide a numerical fitting for the percentage deviation from $\Lambda$CDM of the averaged relative power spectrum $\Delta R = \bar{R}_{k[0.3, 1.0]} -1$ in the $k$-range $[0.3, 1.0]\ h\textrm{Mpc}^{-1}$, as a linear function against $m_{\nu}$ and $\eta^2$. The selection of this $k$-mode range can be customized for different observations. This two-variable regression shows good linearity. The redshift evolution of the corresponding coefficients $C_m$ and $C_{\eta}$ are also studied. We observe that although $C_m$ and $C_{\eta}$ are similar in magnitude in the late time, they differ at higher redshifts. We further investigate the covariance matrix of this numerical fitting and find that the off-diagonal term is comparable to the diagonal terms. Therefore, we propose to characterize the neutrino properties by a redshift independent parameter $m_{\nu} - \frac{4}{3}\eta^2$, which dominates the neutrino effects on the cosmological structure as well as the covariance. This also suggests that LSS alone may not be enough to break the parameter degeneracy between $m_{\nu}$ and $\eta$, and combined analyses with other cosmological probes such as CMB are needed. We choose the grid-based method to include massive neutrinos in our simulation. It is not only more efficient, but can also be easily extended to neutrino models with a non-zero lepton asymmetry parameter $\eta$, which may not be compatible with the particle-based simulation schemes. This is because the non-zero $\eta$ effectively introduces a degeneracy pressure term apart from gravity, which is not a two-body force, and thus difficult to simulate in current N-body schemes. Our code is conceptually similar to that of \cite{ab2013}, only that we have modified \texttt{Gadget2} instead of \texttt{Gadget3}, which makes it much easier to be made public in the future. In this work, we investigate the effects of neutrinos on the matter power spectrum in the comoving units, e.g. $k[h \textrm{Mpc}^{-1}]$, which is the intrinsic unit in cosmological simulations, and thus the systematic error is minimized. However, as the fitting result of $h$ is dependent on $m_\nu$ and $\eta$, we may need to transform to physical units in future comparison with observations. It should also be noted that as mentioned in \cite{barenboim_mass_eigenstate}, by the time of neutrino decoupling, the neutrino lepton asymmetry is already diagonal in mass eigenstates. Thus strictly speaking, cosmological neutrinos do not follow thermal distribution in mass eigenstates, but linear combination of thermal distributions in flavor eigenstates. Further studies into how reliable this assumption can be and how to improve it should be done. The authors of a recent study \cite{inman} comment on the grid-based simulation that the effect of the sound speed is ignored when the neutrinos are treated as fluid, which leads to a deficit in the neutrino power spectrum. This should not affect our results much, as the neutrino power spectrum only accounts for a tiny part in the total matter power spectrum. However, future studies may need to include this correction when high accuracy is required. \appendix | 18 | 8 | 1808.00357 |
1808 | 1808.09871_arXiv.txt | {The study of deuteration in pre-stellar cores is important to understand the physical and chemical initial conditions in the process of star formation. In particular, observations toward pre-stellar cores of methanol and deuterated methanol, solely formed on the surface of dust grains, may provide useful insights on surface processes at low temperatures.}{ Here we analyze maps of CO, methanol, formaldehyde and their deuterated isotopologues toward a well-known pre-stellar core. This study allows us to test current gas-dust chemical models.}{Single-dish observations of CH$_3$OH, CH$_2$DOH, H$_2$CO, H$_2\,^{13}$CO, HDCO, D$_2$CO and C$^{17}$O toward the prototypical pre-stellar core L1544 were performed at the IRAM 30 m telescope. We analyze their column densities, distributions, and compare these observations with gas-grain chemical models. } {The maximum deuterium fraction derived for methanol is [CH$_2$DOH]/[CH$_3$OH] $\sim$ 0.08$\pm$0.02, while the measured deuterium fractions of formaldehyde at the dust peak are [HDCO]/[H$_2$CO] $\sim$ 0.03$\pm$0.02, [D$_2$CO]/[H$_2$CO] $\sim$ 0.04$\pm$0.03 and [D$_2$CO]/[HDCO] $\sim$ 1.2$\pm$0.3. Observations differ significantly from the predictions of models, finding discrepancies between a factor of 10 and a factor of 100 in most cases. It is clear though that to efficiently produce methanol on the surface of dust grains, quantum tunneling diffusion of H atoms must be switched on. It also appears that the currently adopted reactive desorption efficiency of methanol is overestimated and/or that abstraction reactions play an important role. More laboratory work is needed to shed light on the chemistry of methanol, an important precursor of complex organic molecules in space.}{} | Pre-stellar cores form in molecular clouds, due to the influence of gravity, magnetic fields and turbulence. They are dense ($n_{\rm{H}_2}>10^5$\,cm$^{-3}$) and cold ($T<$10\,K) toward their center \citepads{2010MNRAS.402.1625K}. They are the starting point in the process of star formation \citepads{2007ARA&A..45..339B, 2012A&ARv..20...56C}, as they are self-gravitating dense cores which present signs of contraction motions and chemical evolution \citepads{2005ApJ...619..379C}. Hence, they represent ideal laboratories to study the early evolutionary stages of low-mass star formation. Methanol (CH$_3$OH) is one of the most widespread organic molecules in the ISM and a major precursor of chemical complexity in space. Methanol is believed to form on dust grain surfaces by sequential addition of hydrogen atoms to adsorbed CO molecules \citepads{1982A&A...114..245T, 2002ApJ...571L.173W}. This process should take place in cold dense cores at densities above 10$^4$\,cm$^{-3}$, where large amounts of gas-phase CO molecules start to freeze-out onto dust grain surfaces \citepads{1998ApJ...507L.171W, 1999ApJ...523L.165C, 2002ApJ...569..815T}. In L1544, \citetads{2014ApJ...795L...2V} deduced that the CH$_3$OH emission should trace an external layer of the core. In fact, \citetads{2014A&A...569A..27B} mapped the CH$_3$OH emission across L1544, finding an asymmetric ring-like shape surrounding the dust peak. The methanol peak is about 4000 au away from the core center and its distribution can be reproduced by the recent gas-grain chemical model (applied to L1544) by \citetads{2017ApJ...842...33V}. This can be understood by the fact that the CH$_3$OH molecule can more easily desorb from the grain surface when a significant fraction of surface ice is in CO, and CO-rich dust surfaces are present around 4000-5000 au away from the core center because of the fast CO freeze-out at that location and the different CO and H$_2$O photodesorption rates \citepads{2017ApJ...842...33V}. Deuterated methanol should follow the methanol distribution, as they are both formed on the surface of dust grains. However, as already found by previous authors \citepads{2002ApJ...565..344C}, the deuterium fraction increases toward the densest and coldest parts of the core, where CO is mainly in solid form. Toward the center, due to the dissociative recombination of the abundant deuterated forms of H$_3^+$, the D/H atomic ratio increases from the cosmic value \citepads[$\sim$1.5$\times$10$^{-5}$; ]{2006ApJ...647.1106L} to values larger than 0.1 \citepads[e.g.]{2003ApJ...591L..41R}. Thus, toward the core center, D atoms compete with H atoms in the saturation of CO molecules, efficiently producing deuterated methanol \citepads{2012ApJ...760...40A,2012ApJ...748L...3T,2014prpl.conf..859C}. Observations toward L1544 showed that the centroid velocity ($V_{LSR}$) of CH$_2$DOH is about 0.2 km/s lower than that of CH$_3$OH \citepads{2014A&A...569A..27B}; this shift is the same as the one found when comparing the $V_{LSR}$ of C$^{17}$O (1-0), which traces the outer parts of the core, with the $V_{LSR}$ of N$_2$D$^+$ (2-1), mainly tracing the dense central regions where CO is frozen out \citepads{2002ApJ...565..331C}, thus suggesting that CH$_2$DOH is tracing denser regions compared to CH$_3$OH, in agreement with the theoretical expectations. However, this hypothesis is based on the results of single pointing observations, so no information on the spatial distribution of CH$_2$DOH in the core can be inferred. Formaldehyde (H$_2$CO) can be formed both via grain--surface chemistry (in an intermediate step of the formation of methanol) and gas-phase chemistry \citepads[e.g. through reactions of hydrocarbons with oxygen atoms;]{2017iace.book.....Y}. Which route actually dominates is still a matter of debate \citepads{2006A&A...453..949P}. In fact, \citetads{2009A&A...508..737P} suggested that pure gas-phase chemistry can account for the abundances of formaldehyde and its deuterated species observed toward the Orion Bar PDR. On the other hand, \citetads{2011A&A...527A..39B} had to invoke grain--surface chemistry to explain the high deuteration found in the $\rho$ Oph~A cloud. Interestingly, these authors found D$_2$CO to be more abundant than HDCO. The deuterium fraction measured in N$_2$H$^+$, NH$_3$, H$_2$CO and CH$_3$OH \citepads{2009A&A...493...89E, 2017MNRAS.467.3011B} reaches the largest values in dynamically evolved starless cores (pre-stellar cores) and toward the youngest protostellar objects, while it decreases in more evolved phases in the low-mass star formation process. The study of different evolutionary stages is thus important to be able to understand how the deuterium budget builds up in molecules formed preferentially in the gas phase and on dust grain surfaces, and to follow these processes during the dynamical evolution of star forming clouds. We focus our investigation on L1544, a nearby well-known pre-stellar core in the Taurus Molecular Cloud. Here, we study the spatial features of the methanol and formaldehyde deuteration, together with the CO distribution, as CO is thought to be the parent molecule of methanol. The objective is to gain insights on the surface chemistry and the early history of deuteration during the formation of low-mass stars, which is still unknown. The structure of the Chapter is the following: in Section \ref{Observations2} we describe the observations, followed by the results in Section \ref{results2}. The analysis, comparison with models, and discussion are presented in Sections \ref{analysis2}, \ref{models2} and \ref{discussion2}, respectively. Our findings are finally summarized in Section \ref{conclusions2}. | We have presented our maps of methanol, formaldehyde and their deuterated species toward the well known pre-stellar core L1544. These two molecules can help us to gain understanding about the chemical processes taking place on the dust grain surfaces and the formation of more complex organic molecules, as well as the deuterium history in the process of star formation. The highest level of deuteration of methanol occurs close to the dust peak, reaching [CH$_2$DOH]/[CH$_3$OH] = 0.8. This indicates that a more external layer is traced by CH$_2$DOH compared to N$_2$H$^+$ \citepads{2002ApJ...565..331C} and NH$_3$ \citepads{2007A&A...470..221C}. CH$_2$DOH also shows a higher abundance at a distance of $\sim$3000 au from the dust peak, exactly where more methanol is present in the gas phase with respect to CO. This suggests that deuterated methanol is formed and released the same way CH$_3$OH is. HDCO and D$_2$CO, however, peak toward the center of the core, and present a high deuterium fraction, only found previously in $\rho$ Oph A \citepads{2011A&A...527A..39B}. Interestingly, H$_2$CO shows a ring like structure, depleting towards the center, and showing two maxima, one coinciding with C$^{17}$O, and a secondary peak toward the South-West, unlike methanol, which coincides with a region where C$^{17}$O shows less emission. This suggests that gas phase production via reactions involving hydrocarbons efficiently takes place in regions where C is not completely locked in CO, based on the conclusions of \SPEt. Finally, we compared two different chemical models with our observational results, and did not find an agreement. On the one hand, the model from \citetads{2017ApJ...842...33V} is overproducing methanol and formaldehyde, probably caused by the overestimation of the efficiency of reactive desorption. On the other hand, the model from \citetads{Sipila15a,Sipila15b} does not produce enough methanol because of the slow diffusion of hydrogen atoms on the surfaces (which are not allowed to quantum tunnel) and it overproduces formaldehyde. Both models are static and the inclusion of dynamical evolution could change the results, although not by orders of magnitude, needed to reconcile models with observations. Our results suggest that quantum tunnelling for H diffusion on icy dust mantles should be considered, while reactive desorption still needs more detailed experimental work. Higher sensitivity and angular resolution observations of the lines presented here are needed, together with a parameter-space exploration within current chemical models and laboratory work, to shed light on important chemical processes happening at the dawn of star formation. | 18 | 8 | 1808.09871 |
1808 | 1808.05058_arXiv.txt | The Simons Observatory will consist of a single large (6~m diameter) telescope and a number of smaller ($\sim$0.5~m diameter) refracting telescopes designed to measure the polarization of the Cosmic Microwave Background to unprecedented accuracy. The large aperture telescope is the same design as the CCAT-prime telescope, a modified Crossed Dragone design with a field-of-view of over 7.8 degrees diameter at 90~GHz. This paper presents an overview of the cold reimaging optics for this telescope and what drove our choice of 350--400~mm diameter silicon lenses in a 2.4~m cryostat over other possibilities. We will also consider the future expandability of this design to CMB Stage-4 and beyond. | \label{sec:intro} Although the Cosmic Microwave Background (CMB) has already reshaped our knowledge of the contents, origins, and evolution of our universe, the next stage of high precision CMB observations has the potential to revolutionize our understanding. To this aim, the Simons Observatory (SO) Collaboration is building a series of telescopes. Located at an altitude of 5200~m in the Atacama Desert, a site well known for excellent visibilities at millimeter wavelengths, these telescopes will observe the sky with polarization-sensitive detectors operating from 27 to 270~GHz, a frequency range chosen to allow the separation of foreground effects and the polarized CMB signal that contains many different signatures of the early Universe. To achieve these scientific goals, angular scales between a few arcminutes to well over a degree need to be measured with high fidelity -- something a single instrument design would struggle to do. Instead, a large 6~m telescope will be used to measure the small to medium angular scales and a complementary series of $\sim$0.5~m telescopes will be used to recover the larger angular scales. Sufficient overlap in their coverage of different angular scales will allow accurate cross calibration. This paper focuses on the optical design of the camera to be used on the large telescope. First, the design of the telescope is described and next, the cold optics designed to be used with it. An overview of the performance of the chosen cold optical design is given, and possible upgrades to the design are discussed. | We have presented detailed optical designs for the initial deployment on the SO Large telescope. We found 450~mm optics tubes offered the best compromise between sensitivity, modularity and leaves us able to expand in future years. As each optical tube has only 3 lenses and these have identical shapes for all optical tubes we have minimized cost and schedule risk while also enabling parts to be exchanged between optical tubes of the same frequency. Detailed design of lens mounts and baffling is underway and construction should begin shortly. \appendix | 18 | 8 | 1808.05058 |
1808 | 1808.09089_arXiv.txt | We use the semi-analytic model (SAM) of galaxy formation and evolution \sag~coupled with the \textsc{multidark} simulation MDPL2 to study the evolution of the stellar mass-gas metallicity relation of galaxies (MZR). We test several implementations of the dependence of the mass loading due to supernovae (SNe). We find that no evolution in the normalization of the MZR is obtained unless we introduce an explicit scaling of the reheated and ejected mass with redshift as $(1+z)^\beta$. The latter is in agreement with results from the FIRE simulations, and it should encompass small scale properties of the interstellar medium varying over time, which are not captured in SAMs, as well as other energy sources in addition to SNe. Increasing $\beta$ leads to stronger evolution of the MZR normalization; $\beta = 1.9$ reproduces the observed MZR in the range $0 < z < 3.5$. A stronger redshift dependence of outflows reduces the levels of star formation at earlier epochs with the consequent decrease of metal production. This leads to a slower increase of the gas metallicity compared to the stellar mass build-up. The cold gas can be contaminated either by receiving a direct injection of the material recycled by stellar winds and SNe or by gas cooling. The relative role of each process for a given stellar mass depends on the criterion adopted to regulate the fate of the recycled material. However, modifying the metal loading of the outflows has mild impact on the zero-point evolution and does not affect our conclusions. | \label{sec:Intro} The gas-phase oxygen abundance, or gas metallicity, is the measure of the amount of oxygen relative to hydrogen in the gas content of a galaxy. Observations show that metallicity correlates positively with stellar mass \citep{Tremonti2004}. This relation is known as the mass-metallicity relation (MZR) and provides valuable insight onto the impact of the different physical processes that shape the formation and evolution of galaxies. Dying stars produce metals that are returned to the interstellar medium (ISM) via stellar mass loss, resulting in enrichment \citep{Zahid2012}, while gas dilution happens if inflows of pristine gas are present. Thus, these competing mechanisms leave imprints on the oxygen abundance that can be used to infer gas inflows \citep{Zahid2014b,Lagos2016}. Another important physical process that can affect the metallicity are galactic outflows, which can transport part of their metal to the galaxy halo and beyond \citep{AnglesAlcazar2017}.\\ \indent A distinctive characteristic of the MZR is that at stellar masses $\gtrsim 10^{10}$ \msun~the metallicity flattens to an almost constant value \citep{Tremonti2004}. This effect has been attributed to galaxy downsizing \citep{Maiolino2008}, i.e., galaxies with higher stellar masses having their star formation rate (SFR) peaking at higher redshift, in which case the pollution rate of the cold gas in massive galaxies has dropped considerably compared to low mass galaxies, which are still forming stars and injecting new metals onto the ISM. \citet{Zahid2013} argue that the flattening can also be caused by metallicities saturating at an oxygen abundance equivalent to the yield produced by the stellar nucleosynthesis of stars that have enough time to evolve and return all those metals to the ISM. They go even further claiming that the saturation is produced by the equilibrium between the returned metals of dying stars (usually massive ones) and the metals still contained in living stars \citep{Zahid2014a}. Many authors have found that the MZR does evolve in normalization, in a way that high redshift galaxies have a lower oxygen abundance at fixed stellar mass \citep{Erb2006, Maiolino2008, Nagao2008, Mannucci2009, Moller2013, Zahid2014a, Troncoso2014, Wu2016, Ly2016}. It is thought that the evolution of the MZR is a consequence of a complex interaction between the infall of pristine gas towards the galaxies, outflows of metal-rich gas, and mass loss from stars, as it has been suggested from the observations of atomic gas and its relation with metallicity \citep{Bothwell2013,Brown2018}. Another interesting feature of the MZR is the observed scatter, even at $z=0$, as galaxies of the same stellar mass have a wide range of metallicities ($1 \sigma$ of $\pm 0.3$~dex), according to the estimate done by \citet{Mannucci2010}. These authors found that introducing the SFR as an additional variable led to a significant reduction of the scatter to $0.05$~dex, with higher metallicities being associated to lower SFRs at fixed stellar mass. They were guided by the results obtained by \citet{Ellison2007} which show that the MZR has a mild negative correlation with the specific SFR (sSFR) for galaxies with stellar mass $M_\star \lesssim 10^{10}$~\msun. At the same time, \citet{LaraLopez2010} introduce the concept of a Fundamental Plane that correlates the stellar mass, the gas-phase metallicity and the SFR, which seems not to evolve with time. This surface, called Fundamental Metallicity Relation (FMR) by \citet{Mannucci2010}, is of important interest because it could explain the dispersion of the MZR and provide insight into the study of galaxies at high redshift due to its invariance with time. However, the question of the role of the SFR in the MZR is a contingent one. \citet{Sanchez2013} did not find a correlation between metallicity and SFR using integral field spectroscopy, arguing that the FMR was an artefact of the fibre spectroscopy used by previous work, which only sampled the central regions of galaxies, giving an incomplete view of their global metallicity. \citet{Brown2018} showed that depending on the adopted metallicity and SFR calibrations, the correlation between metallicity and SFR can go from being anti- to positively correlated. Observations of atomic (HI) and molecular (H$_2$) hydrogen have been used to suggest that the gas content may be more fundamental than the SFR in determining the scatter of the MZR. On the same line, \citet{Zahid2014a} used an analytical model of chemical evolution to find that the relation between metallicity, stellar mass and hydrogen gas mass does not evolve with time for $z \lesssim 1.6$. Additionally, \citet{Zahid2014b} found observational evidence of the FMR evolution within the same redshift range. Hence, they conclude that the invariance of the FMR found by \citet{Mannucci2010} is attributed to the method used to calculate SFRs, and favour the idea that gas mass is a more fundamental property in the MZR. \citet{Bothwell2013,Hughes2013,Brown2018} have shown that higher metallicities are associated to lower gas fractions at fixed stellar mass. This naturally happens if gas plays the dual role of increasing the gas fraction as well as diluting the metals in the ISM as it inflows \citep{LaraLopez2013a}. To reproduce the MZR locally and at high redshift has been proven to be a very challenging task for simulations. Semi-analytic models (SAMs) such as \textsc{galform} \citep{Cole2000, Bower2006, Gonzalez2014, Lagos2012, Lacey2016} and \textsc{l-galaxies} \citep{Springel2005, Croton2006, Guo2011, Henriques2015} do not produce a MZR that displays the observed normalization evolution \citep{Guo2016}. An explanation for this behaviour might be that the metallicity of galaxies in these SAMs grows at the same rate as the stellar mass, on average. Nonetheless, using the semi-analytic model \textsc{gaea} (\textit{GAlaxy Evolution and Assembly}), \citet{Xie2017} confirmed an evolution of the MZR up to $z \sim 0.7$ for all stellar masses and up to $z \sim 2$ for most massive galaxies (\mstar~$\gtrsim 10^{10}$ M$_\odot$), something that so far could not be achieved by SAMs, which they attributed to a combination of their outflow model and adopted star formation law. Recent cosmological hydrodynamical simulations have had more success reproducing the MZR relation and its evolution. \citet{Dave2011} showed that changing the model of stellar winds in their Smooth Particle Hydrodynamic (SPH) code leads to different rates of MZR evolution due to the impact this has on the interplay between gas inflows and outflows. More recently, \citet{Dave2017} show that their new cosmological hydrodynamic simulation MUFASA displays an MZR that evolves in normalization by $0.4$~dex from $z=0$ to $z=6$. Other cosmological hydrodynamical simulations, such as EAGLE (Evolution and Assembly of GaLaxies and their Environments; \citealt{Crain2015, Schaye2015}) show a similar magnitude of evolution \citep{Guo2016,DeRossi2017}. In addition, \citet{Lagos2016} and \citet{DeRossi2017} showed in EAGLE that the gas metallicity is more fundamentally correlated with gas fraction than SFR. Cosmological zoom-in simulations FIRE (Feedback in Realistic Environments; \citealt{Hopkins2013}) have also been able to reproduce the evolution of the MZR normalization \citep{Ma2016}, while discrepancies with observations become more prevalent in the slope of the MZR. From these results it becomes apparent that hydrodynamical simulations are able to delay the metal enrichment of the ISM of galaxies to a degree that SAMs have struggled with. In this work we tackle this problem and investigate the origin of the evolution of the MZR with the semi-analytic model of galaxy formation and evolution \sag~\citep[][Paper I, hereafter]{Cora2018}, focusing on the impact produced by different modelling of supernova (SN) feedback, and how that changes the interplay between gas inflows/outflows. \sag~features a detailed non-instantaneous recycling chemical enrichment model in which several elements, including oxygen, are followed individually. This makes \sag~an ideal laboratory to perform this study. This paper is organized as follows. In Sec.~\ref{sec:Model}, we give a brief overview of the dark matter (DM) simulation we use to couple with the semi-analytic model \sag~and of the physics included in this model (Paper I), describing with more detail the SN feedback model considered, and how the chemical enrichment of the different baryonic components is calculated. We present the predicted MZR at $z=0$ and at higher redshifts, and compare them with observations in Sec.~\ref{sec:MZR}, showing that the model gives rise to evolution in the normalization of the MZR\footnote{Any time we mention MZR evolution we are referring to the evolution of the normalization of this relation.}. In Sec.~\ref{sec:Analysis}, we analyse the results of two additional versions of the \sag~model that differ in the treatment of the reheated material by SNe with the aim of identifying the physical driver of the MZR evolution found in our reference model. We also demonstrate that the explanation for the origin of the MZR zero-point evolution remains valid when changing the fate of the recycled material. We discuss our results and present our conclusions in Sec.~\ref{sec:Conclusions}. | \label{sec:Conclusions} We have analysed the MZR predictions of the latest version of the semi-analytic model of galaxy formation \sag, described in detail in Paper I. The latest improvements implemented in \sag~include a robust model of environmental effects through the action of RPS and TS coupled to the integration of the orbits of orphan satellites, and a higher efficiency of SN feedback at early times allowed by an explicit redshift dependence of the reheated and ejected mass. The latter is inspired on the results of the hydrodynamical simulation suite FIRE \citep{Muratov2015}. Paper I shows that a variant of the model referred to as \sagcal~allows us to achieve good agreement with observational results for a large range of galaxy properties locally and at high redshift; the SMF at $z=0$ and $z=2$, the SFR function, the scaling relations between black hole and bulge mass, between the cold gas fraction and stellar mass, between the stellar and host halo mass, between the atomic gas content and stellar mass and the evolution of the SFRD. In particular, the detailed non-instantaneous chemical enrichment modelling makes \sag~ideally suited to carry out the present work. Reproducing the development of the MZR has been a long standing problem for SAMs. The importance of the study of the origin of this relation and the way in which it evolves resides in the insights that it provides on the chemical enrichment of the cold and hot gas phases in connection with gas inflows and outflows. We summarize our conclusions as follows: \begin{itemize} \item A SN feedback model which gives rise to larger amounts of reheated and ejected mass at higher redshifts allows us to reproduce the observed evolution of the MZR normalization. This evolution is found for a wide range of stellar masses ($10^{8}$ \msun~$\lesssim$ \mstar~$\lesssim 10^{12}$ \msun) until $z \lesssim 3.5$ for the model \sagcal, in which the new SN feedback recipe involves an explicit redshift dependence (Fig.~\ref{fig:MZRmodel0} and \ref{fig:MZRmodel}). The lack of such dependence, as in the old SN feedback scheme, i.e. the model \oldfb, prevents the MZR normalization from evolving (Fig.~\ref{fig:MZRevolution}). We also find that the new SN feedback scheme yields no evolution of the slope of the MZR. \item We experimented with the redshift dependence of the reheated and ejected mass in the new SN feedback recipe (represented by the parameter $\beta$) and found that a larger $\beta$ directly translates into a stronger evolution of the MZR normalization ($\approx 0.44$~dex). While the \sagbeta~model, which has $\beta = 1.3$, produces some evolution of the MZR normalization ($\approx 0.28$~dex), it is not as pronounced as the one achieved by the \sagcal~model, which adopted $\beta = 1.99$ (Fig.~\ref{fig:MZRevolution}). Higher values of the parameter $\beta$ embody more violent effects of SN feedback at high redshifts. \item Larger feedback effects at higher redshifts produce larger discrepancies between the modelled and observed SFRD, with the former being below recent data by a factor $\approx 3$ (Fig.~\ref{fig:cSFR}). Models with stronger SN feedback effect at higher redshifts produce lower values of SFR, with the consequence of a lower production of metals, as indicated by the cosmic density of oxygen content of the cold gas (Fig.~\ref{fig:Omegagas}). \item Stronger SN feedback at higher redshifts leads to a delay in the pollution of the hot gas. For a given stellar mass, ${\rm Z_{hot}}$ increases with decreasing redshift in the model \sagcal, in contrast with the model \oldfb, in which ${\rm Z_{hot}}$ practically remains the same throughout time at fixed stellar mass (Fig.~\ref{fig:ZHotMstar}). The lower levels of metal enrichment of the hot has halo at higher redshifts in \sagcal~are a result of the lower values of SFR and the larger amounts of reheated mass which dilute the metals injected in that gas phase (Fig.~\ref{fig:ReheatMstar}). The slower increase of the hot gas metallicity compared with the stellar mass build-up is finally translated as an evolution of the MZR normalization, since part of the cold gas contamination takes place through gas cooling. \item In our model, the metal pollution of the cold gas occurs either through direct injection of recycled material into the cold gas or by the gas cooling of enriched hot gas. For a given stellar mass, the relevance of each process depends on the criterion (recycling scheme) adopted to decide the fate of the recycled material, i.e. cold or hot gas phase. When the fraction of recycled mass that ends up in the cold gas phase is defined by the ratio between the reheated and recycled mass (as in the \sagcal, \sagbeta~and \oldfb~models), gas cooling is the main channel of pollution for low-mass galaxies ($M_{\star} \lesssim 10^{10}\,M_{\odot}$) but losses importance as stellar mass increases. The opposite trend is obtained when the fate of the recycled mass is regulated by the cold gas fraction of galaxies (Fig.~\ref{fig:RecycledMstar}). \item The modelling of the metal loading of the outflows does not affect our conclusions. An experiment in which the fate of the recycled material is regulated by the cold gas fraction (model \sagrec, recycling scheme 2), instead of being driven by the amount of reheated mass as in \sagcal~(recycling scheme 1), shows that the galaxy population exhibits an evolution of the MZR normalization that is only slightly smaller ($\approx 0.35$~dex) than our original model (Fig.~\ref{fig:MZRevolution}). This small difference is attributed to a faster enrichment of the cold gas in \sagrec~because of the predominant role of direct metal injection onto this gas phase for low-mass galaxies. \end{itemize} Overall, with SAG we find that a stronger evolution of the SN feedback effect is key in getting the expected evolution of the MZR normalization, which has been a long standing challenge for SAMs. Both the mass of outflows and metals produced are a direct consequence of the action of SN feedback. Thus, the evolution of the zero-point of the MZR is mainly caused by the lower levels of star formation, and consequent lower metal production, as a result of stronger SN feedback at high redshift, with minor dependence on the fate of the metals produced by stellar evolution and returned via stellar winds and SNe. The explicit redshift dependence involved in the estimate of outflows generated by SN feedback is interpreted as evolving aspects of the ISM not captured by the model, such as different redshift evolution of local disc properties with respect to the circular velocity of galaxies \citep{Lagos2013, Creasey2013}, or additional sources of energy (other than SNe) that may play an important role \citep{Hopkins2012}. | 18 | 8 | 1808.09089 |
1808 | 1808.00027_arXiv.txt | Caustic rings of dark matter with \textit{tricusp} cross section were predicted to lie in the galactic disk. Their radii increase on cosmological time scales at a rate of order $1$ kpc/Gyr. When a caustic ring passes through the orbit of a star, the orbit is strongly perturbed. We find that a star moving in a nearly circular orbit is first attracted towards the caustic ring, then moves with and oscillates about the caustic for approximately $1$ Gyr before returning to its original orbit. As a result, a stellar overdensity forms around the caustic ring. We predict such overdensities to be of order $120$\% near the second caustic ring where the Monoceros Ring is observed and of order $45$, $30$ and $15$\% near the third, fourth and fifth caustic rings, respectively. We show that the associated bulk velocities of the stars near the caustic rings are less than a few km/s. We also determine the density profile of interstellar gas near the fifth caustic ring assuming it is in thermal equilibrium in the gravitational potential of the caustic and of the gas itself. | Introduction} From various observations, it is known that $23\%$ of the energy density of the Universe is made of cold dark matter (CDM) \cite{Bertone, *Weinberg, *Kolb}. Axions are one of the widely accepted cold dark matter candidates \cite{Sikivie08, *Marsh16}. The QCD axion was originally proposed as a solution to the strong \textit{CP} problem in the standard model \cite{Peccei77, *Peccei2, *Weinberg78, *Wilczek78, Kim79, *Shifman80, *Zhitnitskii80, *Dine81}. It became a cold dark matter candidate when it was shown that QCD axions are produced abundantly in the early universe with very small velocity dispersion \cite{Preskill83, *Abbott83, *Dine83, *Ipser83}. Since dark matter particles are collisionless, they are described in six-dimensional phase space. Because they have very small velocity dispersion, CDM particles lie on a three-dimensional hypersurface embedded in that space. One inevitable consequence of this is the formation of \textit{caustics} \cite{Zeldovich70, *Arnold82, *Shandarin89} \cite{Sikivie99} \cite{Natarajan05, *Natarajan06}. Caustics are surfaces in physical space where the density is infinite in the limit of zero velocity dispersion. In a galactic halo, both outer and inner caustics are formed. As outer caustics appear in the outer regions of the galactic halo (e.g. at a distance of order $100 \; \text{kpc}$ from the Milky Way center), they hardly affect the stellar dynamics in the disk. In this paper, we only discuss the inner caustics which form when the particles are at their closest approach to the galactic center (see Fig.~\ref{tricusp_flows}). The $n$th inner caustic forms in the flow of particles experiencing their $n$th infall in the galactic potential well. If the total angular momentum of dark matter particles is dominated by net overall rotation, each inner caustic is a closed tube whose cross section has the shape of a tricusp \cite{Sikivie98, Sikivie99} (see Fig.~\ref{tricusp}). This structure is called a \textit{caustic ring} of dark matter. The caustic ring model \cite{Sikivie99, Natarajan06, Duffy08} is a proposal for the full phase space distribution of cold dark matter halos. It is axially symmetric, reflection symmetric and has self-similar time evolution. It predicts that caustic rings lie in the galactic plane and that their radii $a_n(t)$ increase on cosmological time scale as $a_n(t) \propto t^{4/3}$. There is observational evidence in support of the caustic ring model \cite{Sikivie03, Banik17}. Furthermore, it was shown that net overall rotation, self-similarity and axial symmetry are the expected outcomes of the rethermalization of Bose-Einstein condensed axion dark matter \cite{Sikivie09, Erken12} before it falls onto a galactic halo \cite{Sikivie11, Banik13}. In the Milky Way, the present radius of the $n$th caustic ring is approximately $\frac{40 \; \text{kpc}}{n}$. Since the Solar System is about $8.5$ kpc away from the galactic center \cite{Binney}, the $n$ = 1, 2, 3 and 4 caustic rings have passed the solar orbit while the $n = 5$ caustic ring is approaching. We find that star orbits are strongly perturbed by a passing caustic ring (see Fig.~\ref{Eplot}). A star on a nearly circular orbit is first attracted toward the caustic ring, then moves with and oscillates about the caustic ring for approximately $1$ Gyr, and finally returns to its original orbit (see Fig.~\ref{rtplot}). This implies that a star overdensity forms near a caustic ring. The overdensity is determined by the depth $(\Phi_c)$ of gravitational potential well of the caustic and the velocity dispersion $(\sigma)$ of the stellar population near the caustic. If $\sigma^2$ is smaller than $\Phi_c$, the stellar distribution is heavily affected by the caustic. Due to larger infall rates \cite{Duffy08}, caustic rings with smaller $n$, i.e. with larger radii, have stronger gravitational fields. Also, the stars near caustic rings of small $n$, i.e. those in the outer regions of the galactic disk, have relatively small velocity dispersions. As a result, large star overdensities form near the caustic rings of small $n$. We predict such overdensities around various caustic rings (see Figs.~\ref{overdensityX} and \ref{overdensityZ}) by simulating the dynamics of half a million stars for each. We perform the simulations for the radial and vertical motions independently as it is computationally expensive to do so for the coupled two-dimensional motions of a large number of stars. We estimate the total overdensity near a caustic ring as the sum of overdensities formed due to the radial and vertical motions. Large star overdensities would attract more stars and interstellar gas, and are expected to be enhanced further. Such feedbacks are not considered here. The Monoceros Ring \cite{Newberg02, *Yanny03, *Ibata03, *Rocha03} has been observed at the location of the second caustic ring. We find that the star overdensity near the $n=2$ caustic ring is of order $120\%$. This reinforces the claim of Ref.~\cite{Natarajan07} that the Monoceros Ring may be caused by the second caustic ring in our galaxy. We estimate the size of the tricusp of the $n=2$ ring to be $p \sim 2.5$ kpc based on the size of the Monoceros Ring. We find the star overdensity of order $45\%$ near the $n=3$ caustic ring, which may explain the existence of the Binney and Dehnen ring \cite{Binney97} at $13.6$ kpc. The overdensities for caustic rings with larger $n$ are smaller, e.g. approximately $30\%$ and $15 \%$ for $n=4$ and $5$ caustic rings, respectively. Such overdensities may be observed in upcoming astronomical data. Recently, three independent groups \cite{Widrow12, *Carlin13, *Williams13} have observed position dependent bulk velocities of order $10$ km/s for the stars in the extended solar neighborhood. Our work was originally motivated to investigate if the passing of a caustic ring through the solar neighborhood may explain such observations. A caustic ring passing the solar neighborhood moves with speed $\sim 1$ km/s $= 1.02$ kpc/Gyr. In Sec. \ref{IIIBa}, we show that the resultant bulk velocities of the stars are of order $\big( \frac{\Delta d}{d} \big)$ km/s, where $d$ is density of stars and $\frac{\Delta d}{d}$ their relative overdensity near the caustic ring. Even if the overdensity were of order $100 \%$, the bulk velocities are quite small compared to the observed ones. Hence, a passing caustic ring cannot explain the observed bulk velocities. More prominent astrophysical signatures of the caustic rings with large $n$ may be found in the distribution of interstellar gas. The interstellar gas has much smaller velocity dispersion than the stars and is strongly affected by caustic rings. We study the effects of the $n=5$ caustic ring on interstellar gas assuming the gas to be in thermal equilibrium in the gravitational potential of the caustic ring and of the gas itself. The caustic ring is taken to be static here because the dynamics of gas is fast compared to the time scale over which the radius of the caustic ring changes. We find that the density of the gas in a cross-sectional plane of the caustic ring has a triangular shape reminiscent of the tricusp. Triangular features in both tangent directions to the nearest caustic ring ($n=5$) have been observed in the IRAS \cite{Sikivie03, Banik17} and GAIA\cite{Chakrabarty18} maps of the galactic plane. Interestingly, the observed features are sharper than those obtained under the above stated assumptions. A brief outline of the paper is as follows. In Sec. \ref{II}, we describe the caustic ring model and determine the gravitational field and potential of a caustic ring. In Sec. \ref{III}, we study the dynamics of the stars in the vicinity of a caustic ring. In sec. \ref{IV}, we study its effect on the distribution of interstellar gas. Sec. \ref{V} provides a summary. | 18 | 8 | 1808.00027 |
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1808 | 1808.02508_arXiv.txt | Short-orbital period small-eccentricity binary pulsars can, in principle, experience substantial advance of periastron. We explore the possibility of measuring this effect by implementing a timing model, {\ttfamily ELL1k}, in the popular {\ttfamily TEMPO2} pulsar timing package. True secular variations in the Laplace-Lagrange parameters, present in our {\ttfamily ELL1k} model, can lead to measurable timing residuals while pursuing decade-long timing campaigns using the existing {\ttfamily ELL1} timing model \good{of \citet{Lange2001}}, especially for binaries exhibiting significant periastron advance. We also list {the} \good{main} differences between our approach \good{and various implementations of the \texttt{ELL1} model present in \good{both} \texttt{TEMPO} and \texttt{TEMPO2}} packages. Detailed {\ttfamily TEMPO2} simulations suggest the possibility of constraining the apsidal motion constant of pulsar companions in certain observed binary pulsars with minuscule eccentricities such as PSR J1719$-$1438. Fortunately, the {\ttfamily ELL1k} timing model does not pose any challenges to the on-going Pulsar Timing Array campaigns that routinely employ the {\ttfamily ELL1} timing model. % | \label{sec:Intro} Presently, binary pulsars provide the most accurate laboratories to test general relativity (GR) in quasi-stationary strong field regime \citep{Wex2014}. This is mainly because of the technique of pulsar timing which requires an accurate prescription for determining the pulse phase as a function of time. In the case of binary pulsars like PSR B1913+16, this technique essentially provided an ideal clock for probing the nature of relativistic gravity \citep{Taylor1979}. Pulsar timing demonstrated the decay of orbital period in neutron star binaries which provided the first observational evidence for the existence of gravitational waves \citep{Tayor_Nobel_lecture}. Additionally, the on-going Pulsar Timing Array (PTA) experiments aim to detect low-frequency gravitational waves in the $10^{-9} \,-\, 10^{-10}$ Hz frequency range \citep{vlh+16}. These efforts require accurate timing of millisecond pulsars (MSPs) due to their exquisite rotational stability. Roughly, 10\% of the \good{over 2600 currently known} pulsars exist in binary systems with companions in various stages of stellar evolution \citep{Manchester1993_ATNF_Catalog}. Accurate timing of pulsars in such systems requires prescriptions to model delays in the times of arrival (TOAs) of pulses in an inertial barycentric frame due to the orbital motion of the pulsar and its companion \citep{Blandford1976,Edwards2006_TEMPO2}. For double neutron star binaries in eccentric orbits, a number of relativistic orbital effects contribute to such delays \citep{DamourDeruelle1986}. This ensures that the binary orbit is specified by a number of post-Keplerian parameters in addition to the regular Keplerian parameters, namely, {the orbital period} $P_a$, {argument of periastron} $\omega$, {eccentricity} $e$, {time of periastron passage} $T_0$ and {the projected semi-major axis} \good{$x$}. Timing of binary pulsars allow us to probe quasi-stationary strong gravitational fields due to our ability to measure their post-Keplerian parameters. However, timing of MSPs in binaries requires special care as a good fraction of them are in near-circular orbits. For pulsar binaries with tiny orbital eccentricities, the TOAs do not prominently depend on Keplerian parameters $\omega$ and $T_0$. This results in large uncertainties in the usual $\chi^2$ estimation of $\omega$ and $T_0$. This prompted Norbert Wex to describe such orbits in terms of the first and second Laplace-Lagrange (LL) parameters and certain time of ascending node passage \good{$T_{\ascnode}$} \citep{Lange2001,Edwards2006_TEMPO2}. These parameters, namely $\epsilon_1 = e\, \sin \omega$, $\epsilon_2 = e\, \cos \omega$ and $T_{\ascnode} = T_{0} - \frac{\omega}{n}$, where $n$ is the mean motion, replace the regular Keplerian parameters $e$, $\omega$ and $T_0$. Clearly, the use of rectangular components of the eccentricity vector, % given by $e\,\cos \omega$ and $e\,\sin \omega$, to represent the periastron of the elliptical orbit is influenced by its use in celestial mechanics \citep{Pannekoek1948}. The {\ttfamily ELL1} timing model, detailed in \citet{Lange2001}, incorporates the effects of \Romer and Shapiro delays while neglecting the Einstein delay contributions. Additionally, {linear-in-time} variations of the sidereal orbital period, Laplace-Lagrange parameters and projected semi-major axis were introduced to model secular changes in these parameters. In the {\tempo} implementation of this timing model, one employs ${\epsilon_{10}}$, $\epsilon_{20}$, $\dot \epsilon_1$, $\dot \epsilon_2$ and $T_{\ascnode}$ as fitting parameters in the place of $e_0$, $\omega_0$, $\dot e$, $\dot \omega$ and $T_0$ employed in the \texttt{DD} timing model for eccentric binaries \citep{DamourDeruelle1986,Edwards2006_TEMPO2}. This is what one gathers from a detailed study of the appendix in \citet{Lange2001} and its {\tempo} implementation \citep{Edwards2006_TEMPO2}. It should be noted that the {\ttfamily ELL1} model incorporates only the first order terms in orbital eccentricity. Very recently, it was pointed out by \citet{Zhu2018} that higher order eccentricity contributions to the \Romer delay should be relevant while timing nearly-circular wide-orbit binary pulsars. \good{These corrections were implemented in the \texttt{ELL1+} model, an extension of \texttt{ELL1}.} However, TOA measurement errors are expected to be much larger than the higher order $e$ corrections to the \Romer delay for short orbital period binaries: systems of our current interest. In this paper, we explore the implications of restricting the temporal variations of LL parameters to be linear-in-time, pursued in \citet{Lange2001} and implemented in {\tempo} \citep{Edwards2006_TEMPO2}. It turns out that the time evolutions of rectangular components of the orbital eccentricity vector are more general. This is because the rate of periastron advance per orbit need not be negligible for compact pulsar binaries with tiny orbital eccentricities as evident from equations (1-3) in \citet{Willems2008}. The more general time evolutions for the LL parameters can have observational implications due to the following reasons. These days it is indeed possible to measure TOAs with \good{$\sim$}100 ns uncertainties due to the advent of real-time coherent de-dispersion \citep{Hankins1987} and availability of large collecting area telescopes like the Giant Metre-wave Radio Telescope equipped with 400 MHz wide-band receivers \citep{Reddy2017}. The TOA precision is likely to improve further with the advent of telescopes like the Five hundred meter Aperture Spherical Telescope (FAST; \citet{Nan2011}) and the Square Kilometre Array (SKA; \citet{Combes2015}). Secondly, the secular systematics in timing residuals are expected to increase with observation span for timing campaigns lasting decades, such as those employed in various PTA experiments \citep{vlh+16}. This is relevant for us as many PTA millisecond pulsars are in small-eccentricity binary orbits and are timed by the {\ttfamily ELL1} timing model. % The final consideration is based on the possibility that the present {investigation} may allow us to extract additional astrophysical information from certain binary pulsar systems like PSR J1719$-$1438 \citep{Bailes2011}. This is because the classical contributions to the periastron advance may dominate over its general relativistic counterpart in a number of recently discovered MSP binaries. If this is indeed the case, the long-term timing of such tiny-eccentricity pulsar binaries may allow us to determine the apsidal motion constant of pulsar companions in such systems. After completion of our numerical {\tempo} experiments, we came to know that Norbert Wex had implemented a similar extension to his {\ttfamily ELL1} timing model in {\texttt{TEMPO}} \citep{tempo1_ref}. Unfortunately, no public documentation exists for this implementation. A close inspection reveals that our timing model and Wex's {\texttt{TEMPO}} implementation\footnote{Available at \url{https://sourceforge.net/p/tempo/tempo/ci/master/tree/src/bnryell1.f}} that extends his {\ttfamily ELL1} timing model, detailed in \citet{Lange2001}, differs by a term. Our {\tempo} simulations reveal that this additional term is crucial for small-eccentricity binary pulsars experiencing substantial periastron advance and therefore should not be ignored. In % \good{Section \ref{sec:RomerDelay_ELL1}}, we briefly summarize the {\ttfamily ELL1} timing model. How we incorporate `exact' temporal evolutions of the LL parameters and why long-term monitoring will be required to distinguish such evolutions from the present {\tempo} evolutions for $\epsilon_1$ and $\epsilon_2$ are explained in Section \ref{sec:Advance-of-periapsis_ELL1}. Why PSR J1719$-$1438 should be the most promising source to observe the effects of periastron advance is explained in Section \ref{sec:measurability}. This section also details how we adapted the {\tempo} software\footnote{Available at \url{https://bitbucket.org/psrsoft/tempo2}} \citep{Hobbs2006_TEMPO2,Edwards2006_TEMPO2} to explore observational consequences of the general time evolutions for $\epsilon_1$ and $\epsilon_2$. In Section \ref{sec:discussion}, we summarize our results and discuss their implications. Expressions required to incorporate the {\ttfamily ELL1k} model in {\tempo} are given in Appendix \ref{sec:t2_inputs} \good{and a brief comparison of different timing models for near-circular binaries is given in Appendix \ref{sec:ell1_comparison_table}.} | \label{sec:discussion} We explored the possibility of measuring the effect of periastron advance in small-eccentricity binary pulsars. This was pursued by implementing a timing model, {\ttfamily ELL1k}, in {\tempo} that essentially incorporated general temporal evolutions for the crucial LL parameters present in the {\ttfamily ELL1} timing model. This present {prescription} should be relevant for binary pulsars that can, in principle, experience significant periastron advance. With the help of {\tempo} simulations and using PSR J1719$-$1438 as an example, we probed the observational implications of {the} {\ttfamily ELL1k} timing model in comparison with the {\tempo} implementation of Wex's {\ttfamily ELL1} model. This comparison provides significant timing residuals while simulating a decade-long timing of PSR J1719$-$1438 even while using the conservative $\dot{\omega}_{\text{GR}}$ estimate for the periastron advance. Our {\tempo} simulations suggest the possibility of measuring the effect of $\dot\omega$ in PSR J1719$-$1438 by pursuing a high-cadence decade-long timing campaign. Additionally, we investigated the measurability of periastron advance in the presence of red timing noise in our {\tempo} simulations for PSR J1719$-$1438. This is achieved by injecting several instances of red noise with varying spectral properties into the simulated TOAs. We concluded that $\dot\omega$ can be measured to about 1\% accuracy provided the amplitude of the red noise is sufficiently small. Furthermore, the tolerable amplitude of red-noise increases with the red-noise timescale ($1/f_c$). Interestingly, if the measured $\dot\omega$ turns out to be larger than the expected $\dot\omega_\text{GR} \sim$ 13$\degree$/yr, it should lead to constraints on the apsidal motion constant $k_2$ of its unique companion. We also compiled a list of binary pulsars with tiny orbital eccentricities where $\dot\omega$ effects may become noticeable within a reasonable timescale especially if the classical contributions to periastron advance dominates over their GR counterpart. More importantly, we showed that the linear-in-time evolutions of LL parameters do not introduce any significant timing residuals in the currently employed PTA pulsars with tiny orbital eccentricities. The present effort should be interesting for the following scenarios. The first one involves the detection of sub-$\mu$Hz GWs using an ensemble of MSPs in a PTA experiment that demands post-fit timing residuals of the order of 10 nano-seconds. Unfortunately, this is not attainable for most of the current PTA MSPs and many of them are part of nearly circular binary systems. A possible source of higher post-fit timing residuals can be unmodeled systematic effects as evident from our plots. Therefore, the use of a more refined timing model can, in principle, lower post-fit residuals, especially for MSPs in compact orbits with small eccentricities. However, it is unlikely that {the} {\ttfamily ELL1k} model will lead to any improvements to the current PTA sensitivity and this is mainly due to the large $\tau_{{\dot \omega}}$ values for the current list of PTA pulsars. However, it will be helpful to check the relevance of {this} timing model while including new MSPs into the existing list of PTA pulsars. The second scenario involves accreting millisecond X-ray pulsars like SAX J1808.4$-$3658. It should be interesting to explore the implications of {this} model while analyzing timing data associated with its many observed outbursts. For such accretion-powered systems, the classical contributions to $\dot \omega$ can be quite large and therefore it may be worthwhile to perform coherent timing of its outburst data by employing our timing model. This may, in principle, lead to an estimate for the $k_2$ value of its brown-dwarf companion and detailed analysis and simulations will be required to quantify our suggestion. Note that it will require LISA observations of eccentric galactic binaries for estimating $k_2$ values of degenerate objects \citep{Willems2008}. Finally, {this} timing model should be relevant during the FAST-SKA era as the MSP population is expected to quadruple during this era \citep{Levin2017}. It is reasonable to expect that SKA will monitor dozens of short orbital period MSPs % in nearly circular orbits and {the} {\ttfamily ELL1k} timing model will be required for the high-cadence timing of such systems. | 18 | 8 | 1808.02508 |
1808 | 1808.03097_arXiv.txt | Flares produce sudden and permanent changes in the horizontal photospheric magnetic field. In particular flares generally produce increased magnetic shear in the photospheric field along the neutral line. Recent observations show also that flares can produce sudden photospheric motion. We present a model for the observed changes as the response of the photosphere to a large-amplitude shear Alfv\'{e}n wave propagating down from the corona on either side of the neutral line. The Alfv\'{e}nic front is assumed to impact the photosphere close to the neutral line first, and then successively further away with time, such that the line of impact coincides with the flare ribbon. The wave introduces magnetic shear and velocity shear. The magnetic shear introduced at the photosphere has the same sign on either side of the neutral line, while the velocity shear has the opposite sign. We discuss the possibility that this process is responsible for particle acceleration in flares. | \label{sec:introduction} During the impulsive phase of a solar flare, magnetic energy is converted into other forms of energy in the solar corona. The accepted mechanism underlying flares is magnetic reconnection, a process in which coronal magnetic field lines change their connectivity. Despite decades of investigation, many details of the flare process remain poorly understood~\citep{2017LRSP...14....2B}. It is difficult to measure the coronal magnetic field directly, but photospheric magnetic field measurements have revealed that flares produce sudden and permanent changes in the observed magnetic field~\citep[e.g.][]{2005ApJ...635..647S}. The most detailed information comes from vector magnetogram measurements, which show that the predominant change in a flare occurs in the horizontal magnetic field, which tends to increase parallel to the neutral line, i.e.\ the magnetic shear along the neutral line increases~\citep{2012ApJ...745L..17W,2012ApJ...759...50P}. There are corresponding sudden changes in the electric current density close to the neutral line~\citep[e.g.][]{2016A&A...591A.141J}. The observations have been interpreted as the photospheric response to coronal magnetic restructuring during the flare. Changes in the photospheric field values imply changes in the net Lorentz force on the corona (which can be calculated from the boundary values of the field), and the values of the changes have been used to interpret the effect of the flare in the corona~\citep{2012SoPh..277...59F,2016SoPh..291..791P,2016ApJ...820L..21X}. \citet{2016NatCo...713104L} reported a striking example of changes at the photosphere during the 22 June 2016 M6.5 flare: a sunspot was observed to rotate suddenly in response to the passage of a flare ribbon across the spot. The observations confirm that coronal field changes can produce not only photospheric field changes, but also substantial induced motion of the dense photosphere, contrary to general expectations~\citep{2016NatPh..12..998A}. Other examples of flare-induced sunspot rotation have also been reported, including the case of a sunspot reversing its direction of rotation~\citep{2016NatCo...713798B}. Hard X-ray (HXR) observations of flares imply that a significant fraction of the released energy goes into accelerated electrons with energies $10-100$\,keV~\citep{2017LRSP...14....2B}. It is generally assumed that the electrons originate high in the solar corona, perhaps at the site of magnetic reconnection, and then follow field lines down to the dense lower atmosphere, where they produce hard X-rays via thick-target bremsstrahlung~\citep{1971SoPh...18..489B}. However, this picture for HXR production suffers from the ``number problem.'' Because the electrons originate in the low-density corona, the implied particle fluxes at the low atmosphere would evacuate electrons from a substantial volume above an active region during a flare~\citep{1976RSPTA.281..473B}. A return current of electrons from the dense chromosphere to the corona is required, but this also introduces problems. The observations of flare-induced photospheric motion imply that energy is also transported from the corona to the photosphere in other forms. The photospheric changes occur behind the flare ribbons, the site of hard X-ray emission, which suggests a more direct connection between the magnetic field change at the photosphere and the acceleration process. In this article we present a simple 2D magneto-hydrodynamic model for the response of the photosphere to a flare, in terms of a large-amplitude shear Alfv\'{e}n wave produced by coronal reconfiguration impacting the photosphere, and introducing a magnetic and velocity shear close to the neutral line. To motivate the model we return to the observations of the 22 June 2015 event (Section~\ref{sec:analysis}). We present a summary analysis of the observations, as well as the results of nonlinear force-free modeling. In Section~\ref{sec:model} we give the details of the model, and in Section~\ref{sec:discussion} we discuss the model results, and a possible connection to electron acceleration in flares. In Section~\ref{sec:conclusions} we draw conclusions. | \label{sec:conclusions} We present a 2-D model which explains the sudden appearance of magnetic and velocity shear at the photosphere during a flare in terms of a downwards-directed large amplitude Alfv\'{e}n shear wave impacting the photosphere on either side of the neutral line. The shear Alfv\'{e}n wave is assumed to be produced by magnetic field reconnection in the flare. Although the wave propagates vertically downwards, the front is assumed to be inclined to the photosphere, so that the front arrives first at locations closer to the neutral line. This is intended to reproduce the observations of a sudden photospheric response to a flare behind spreading flare ribbons~\citep{2016NatCo...713104L}. The model front is transmitted and reflected at the photosphere, and the transmitted wave introduces a horizontal component in the magnetic field, and a horizontal flow, beneath the photosphere. In principle this can account for the surprising observations of sudden motion of the photosphere in response to a flare~\citep{2016NatPh..12..998A}. The model predicts that the shear introduced by the wave in the photospheric magnetic field has the same sign on either side of the neutral line, whereas the velocity shear has the opposite sign. Also, the total energy deposited in the photosphere by the shear Alfv\'{e}n wave is comparable to the flare energy. We speculate that the sudden changes in the magnetic field in the low atmosphere are associated with particle acceleration in the flare. The model is highly simplified, but in principle it can account for a range of effects due to a flare. It remains to work out the details, and to develop more detailed models. | 18 | 8 | 1808.03097 |
1808 | 1808.00577_arXiv.txt | Evolutionary studies that compare galaxy structure as a function of redshift are complicated by the fact that any particular galaxy's appearance depends in part on the rest-frame wavelength of the observation. This leads to the necessity for a "morphological k-correction" between different pass-bands, especially when comparing the rest-frame optical or infrared (IR) to the ultraviolet (UV). This is of particular concern for high redshift studies that are conducted in the rest-frame UV. We investigate the effects of this "band-pass shifting" out of the UV by quantifying nearby galaxy structure via "CAS parameters" (concentration, asymmetry, and clumpiness). For this study we combine pan-chromatic data from the UV through the near-IR with GALEX (Galaxy Evolution Explorer) data of 2073 nearby galaxies in the "NUV" ($\sim 230$ nm) and 1127 in the "FUV" ($\sim 150$ nm), providing the largest study of this kind in the mid- to far-UV. We find a relationship between the CAS parameters and observed rest-frame wavelength that make galaxies appear more late-type at shorter wavelengths, particularly in the UV. The effect is strongest for E/S0 galaxies in the far-UV, which have concentrations and asymmetries that more resemble those of spiral and peculiar/merging galaxies in the optical. This may be explained by extended disks containing recent star formation. Here we also release the CAS values of the galaxies imaged in GALEX NUV and FUV for use in comparisons with deep HST imaging and JWST in the future. | The majority of high redshift (z $> 2$) galaxies appear similar to a relatively rare subset of low redshift irregular and peculiar galaxies, whose morphologies appear to be pathological due to mergers or interactions (e.g., Driver et al. 1995, Conselice et al. 2005, Kriek et al. 2009, Delgado-Serrano et al. 2010, Mortlock et al. 2013). This observed increase in the percentage of apparently merging/interacting galaxies with redshift supports models of hierarchical galaxy formation (e.g. Nagashima et al. 2002). Comparison studies of galaxies at different redshifts are complicated, however, by the fact that galaxies can appear substantially different at shorter wavelengths than at longer ones (e.g., Bohlin et al. 1991, Kuchinski et al. 2000, Windhorst et al. 2002). This is especially true for galaxies that appear as earlier types at the longer wavelengths. This leads to a "morphological k-correction" for a given galaxy between different rest-frame wavelengths. This is particularly important in studies of high redshift galaxies, as band-pass shifting will cause light originally emitted in the UV to be shifted as far as the IR by the time it reaches Earth. This raises questions about how much of the irregular/peculiar morphologies seen in high redshift studies are simply due to band-pass shifting, and not due to real differences in galaxy type. Kuchinski et al. (2001) showed that this effect alone is enough to misclassify galaxies as peculiar when simulated at higher redshifts. A quantitative measurement of the morphological k-correction is therefore an essential first step to the study of galaxy evolution. This correction is expected to be particularly significant when comparing the rest-frame UV to the optical or IR, as galaxy stellar energy distributions (SED's) can change drastically short-ward of the Balmer Break ($\lambda_c \lesssim 360$~nm), and UV-bright star-forming regions dominate galaxy morphology that appears to be smoother at redder wavelengths. To determine the extent of this effect, it is critical to conduct a large representative benchmark study at $z \simeq 0$. This is complicated by the difficulty of observing in the UV through the atmosphere: it is necessary to use space-based telescopes with UV sensitive detectors, such as HST (Hubble Space Telescope) and GALEX (Galaxy Evolution Explorer). Throughout its operation, GALEX has amassed the largest image database in the far-UV to date (Bianchi \& GALEX team 1999, Martin \& GALEX team 2005, 2009). This makes GALEX ideal for these types of studies, where large number statistics are needed to determine overall generalized trends in galaxy structure. One of the most efficient ways to quantify the morphological k-correction is through measurements of the concentration, asymmetry, and clumpiness (CAS) parameters of galaxies as a function of their rest-frame wavelength. The concentration of light as it is radially distributed within a galaxy (C) correlates with the stellar and kinematic masses of galaxies. The degree to which a galaxy deviates from perfect symmetry (A) can indicate on-going galaxy mergers or interactions. The proportion of high frequency structure to the smooth light distribution of a galaxy (S) correlates with current star formation (Conselice 2003). Conselice (2006) argues that these parameters are the more fundamental properties that describe a galaxy's physical state, and thus are a good objective way of classifying galaxies. Nearby galaxy studies have shown that a combination of these parameters provides a relatively robust automated galaxy classification system, which can be used to determine the nature of the population distributions for large samples of galaxies (e.g. Bershady et al. 2000, Conselice 2003, Conselice 2006). In Taylor-Mager et al. (2007), we found that earlier-type galaxies are on average increasingly more concentrated, more symmetric, and less clumpy than later-type galaxies. Merging galaxies occupy separate locations within the CAS parameter space, which leads to a clear and efficient method of automatically identifying merging galaxies. There have been a few studies that quantify the morphological k-correction of nearby (z $\sim 0$) galaxies (e.g., Bohlin et al. 1991, Windhorst et al. 2002, Vika et al. 2013), one of the largest of which is our earlier project (Taylor-Mager et al. 2007) that utilized multi-wavelength (0.15-0.85$\mu$m) CAS measurements of 199 galaxies. Only a subset of this sample had UV data at the time, with GALEX observations for just 14 of them. Due to signal-to-noise (S/N) and resolution constraints, only 7 galaxies were used to calculate the final k-correction in the NUV ($\lambda_c=2275$\AA) and 5 in the FUV ($\lambda_c=1550$\AA). Complimentary HST WFPC2 data allowed for measurements at longer UV wavelengths, with 8 galaxies in the F255W ($\lambda_c=2550$\AA) filter and 86 in F300W ($\lambda_c=2930$\AA). Galaxy number statistics were small in the UV for all galaxy types, but particularly for early-type (E--S0) galaxies, which were not the focus of that study, and which are especially faint in the UV due to their intrinsically red colors. As such, we were unable to provide a reliable quantified measure of the morphological k-correction in the UV for E--S0 galaxies in that paper. We noted an intriguing trend of E--S0 galaxies apparently becoming dramatically less concentrated in the UV, but included the caveat that this needed to be verified with a larger sample of early-type galaxies. This has motivated our present study, which uses images in the GALEX archive to drastically improve our number statistics (by up to almost a factor of 300) in the far- to mid-UV, providing the largest analysis of the wavelength-dependence of nearby galaxy CAS structure in this wavelength range to date. | In conclusion, we have measured the concentration, asymmetry, and clumpiness (CAS) parameters of 2073 nearby galaxies imaged in the far-UV by GALEX, and release those values for use in future comparison studies, such as those with deep HST and JWST images. We find a significant morphological k-correction in C for E/S0 and Sa-Sc galaxies, with decreasing concentration at shorter wavelengths. There is a particularly dramatic difference in concentration for E/S0 galaxies between the far- and mid-UV. Sd-Im and peculiar/merging galaxies show a weaker trend in C, where concentration does not vary much between filters. There is a strong morphological k-correction in A for all galaxy types, with increasing asymmetry at shorter wavelengths. However, there is little change in A for E/S0 galaxies within optical and IR wavelengths. There is also a strong k-correction in S of increasing clumpiness at shorter wavelengths for spirals, irregulars, and mergers/peculiars. For these galaxy types, we find that care should be taken when comparing images of vastly different spatial resolutions, as much lower resolution leads to a much lower clumpiness value. This resolution effect is not as strong for the inherently smooth E/S0, for which we find little dependence of S on wavelength, except for a mild increase in clumpiness when transitioning from optical to UV wavelengths. The A values, on the other hand, are not significantly affected by image resolution, which show that asymmetries are large scale features that do not disappear at the lower (5 arsec stellar FWHM) GALEX resolution. The relationship between the CAS parameters and the observed rest-frame wavelength makes galaxies in general appear more late-type than they really are at shorter wavelengths, especially in the mid- to far-UV, where the morphologies with CAS become nearly degenerate for all but the early type galaxies (as apparent in Figures 3 and 4). Surprisingly, this effect is so strong for E/S0 galaxies in the far-UV that their concentrations and asymmetries more closely resemble those of spirals and peculiar/merging galaxies at red optical wavelengths. These high asymmetries and low concentrations of ellipticals and lenticulars at shorter wavelengths can be explained by extended ultraviolet disks and halos that have been found in many spiral and irregular galaxies, where low-level star formation is ongoing in the faint outskirts (e.g. Holwerda et al. 2012, Hodges-Kluck et al. 2016). Most surprisingly, early-type (E/S0) galaxies historically considered to be "red and dead" have been found to be UV-bright (Schawinski et al. 2007), many with extended disks visible primarily in the UV (Rutkowski et al. 2012, Petty et al. 2013). This accounts for their diffuse and asymmetric structure as measured here in the UV, resulting in morphologies that more resemble those of disk galaxies in the optical. This UV excess in early-type galaxies can potentially be explained by recent star formation, or an evolved horizontal branch stellar population (O'Connell 1999; "UV upturn"). The stellar population in the inner regions of some E/S0 galaxies have been found to be older than the outer regions (Petty et al. 2013), with more recent star formation in extended H I-rich regions (Yildiz et al. 2017). In at least some cases this may be the result of recent interactions with companion galaxies or gas-rich merger events (Koshy \& Zingade 2015). | 18 | 8 | 1808.00577 |
1808 | 1808.05955_arXiv.txt | {We forecast the sensitivity of thirty-five different combinations of future Cosmic Microwave Background and Large Scale Structure data sets to cosmological parameters and to the total neutrino mass. We work under conservative assumptions accounting for uncertainties in the modelling of systematics. In particular, for galaxy redshift surveys, we remove the information coming from non-linear scales. We use Bayesian parameter extraction from mock likelihoods to avoid Fisher matrix uncertainties. Our grid of results allows for a direct comparison between the sensitivity of different data sets. We find that future surveys will measure the neutrino mass with high significance and will not be substantially affected by potential parameter degeneracies between neutrino masses, the density of relativistic relics, and a possible time-varying equation of state of Dark Energy.} \begin{document} \hfill{\small TTK-18-29} | 18 | 8 | 1808.05955 |
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1808 | 1808.07482_arXiv.txt | Eclipsing binaries in star clusters offer more stringent tests of stellar evolution theory than field binaries because models must not only match the binary properties, but also the radiative properties of all other cluster members at a single chemical composition and a single age. Here we report new spectroscopic observations of the G type, detached eclipsing binary \epic\ in the open cluster Ruprecht~147 (${\rm [Fe/H]} = +0.10$), which was observed in late 2015 by the \ktwo\ mission. A joint analysis of our radial-velocity measurements and the \ktwo\ light curve shows the 6.5 day orbit to be very nearly circular. We derive highly precise masses of $1.0782^{+0.0019}_{-0.0019}$~$\mathcal{M}_{\sun}^{\rm N}$ and $1.0661^{+0.0027}_{-0.0021}$~$\mathcal{M}_{\sun}^{\rm N}$, radii of $1.055 \pm 0.011$~$\mathcal{R}_{\sun}^{\rm N}$ and $1.042 \pm 0.012$~$\mathcal{R}_{\sun}^{\rm N}$, and effective temperatures of 5930~$\pm$~100~K and 5880~$\pm$~100~K for the primary and secondary, respectively. The distance we infer, $283^{+18}_{-16}$~pc, corresponds to a parallax in good agreement with the {\it Gaia}/DR2 value for the star. Current stellar evolution models from the MIST and PARSEC series match the above physical properties very well at ages of 2.48 and 2.65~Gyr. Isochrones for these same ages and the measured composition, along with our reddening estimate for \epic, also show generally good agreement with the optical and near-infrared color-magnitude diagrams of the cluster, which can be constructed with no free parameters as the distances of all member stars are known from {\it Gaia}. | \label{sec:introduction} Star clusters have long been used to test our knowledge of stellar physics, by comparing observations in the color-magnitude diagram against models of stellar evolution. This allows inferences to be made about the age of the cluster and its distance, often aided by knowledge of the chemical abundance that can be derived spectroscopically from one or more of its brighter members. Constraints of a very different nature on stellar theory may be obtained from suitable detached double-lined eclipsing binaries by measuring their component masses and radii \citep[see, e.g.,][]{Andersen:1991, Torres:2010a}, two of the most fundamental stellar properties. Accurate values for these properties can usually be derived from purely geometric and dynamical principles, with no further assumptions. When eclipsing binaries are found in a cluster of known metallicity the constraints become much stronger and the test more valuable, as models must then match not only the masses and radii of both binary components at the measured composition, but also the radiative properties of all other cluster members at the same age as inferred for the binary. While some three dozen eclipsing binaries have been studied photometrically and spectroscopically in clusters and young associations, not all binaries are suitable (detached) for this sort of test, or have parent populations sufficiently well characterized (well defined color-magnitude diagrams, spectroscopically known metallicity), or have had their properties measured well enough. Some examples of systems with sufficiently precise measurements permitting such tests include those published by \cite{Kaluzny:2006}, \cite{Meibom:2009}, \cite{Brogaard:2011, Brogaard:2012}, \cite{Sandquist:2016}, \cite{Brewer:2016}, and others. Here we present an analysis of a relatively bright ($V = 11.4$) detached eclipsing binary in the open cluster Ruprecht~147 (NGC~6774), designated \epic\ (also TYC~6296-96-1 and 2MASS~J19152465-1651222). Ruprecht~147 lies on the ecliptic, and was observed by \ktwo, the re-purposed \kepler\ mission, in the final months of 2015. \epic\ was identified as a cluster member and a double-lined spectroscopic binary by \cite{Curtis:2013}, and listed as entry CWW~64 in their catalog. These authors presented the first detailed study of this neglected cluster, establishing it to be middle-aged ($\sim$3~Gyr), of slightly super-solar composition (${\rm [Fe/H]} = +0.10$), and distant some 300~pc from the Sun, making it the oldest nearby cluster (see their Figure~1). What makes Ruprecht~147 special is that it contains no fewer than \emph{five} eclipsing binaries \citep{Curtis:2016}, offering an unprecedented opportunity for testing stellar evolution models over a wide range of masses at a single age and composition. Furthermore, the recently published second data release (DR2) from the {\it Gaia} consortium \citep{Gaia:2016, Gaia:2018} now provides highly accurate parallaxes for other cluster members that greatly simplifies the study of the color-magnitude diagram of Ruprecht~147. This paper begins our study of the eclipsing systems in the cluster with \epic, a 6.527~day detached eclipsing binary composed of similar G-type main-sequence stars. We describe our preparation of the \ktwo\ data for analysis in Section~\ref{sec:photometry}, our imaging observations to examine the vicinity of \epic\ in Section~\ref{sec:imaging}, and our spectroscopic monitoring to derive radial velocities for both components in Section~\ref{sec:spectroscopy}. The joint analysis of the \ktwo\ light curve and the velocities is presented in Section~\ref{sec:lightcurve}. We then proceed to derive the absolute dimensions of the components (masses, radii, etc.)\ in Section~\ref{sec:dimensions}. Rotation is studied in Section~\ref{sec:rotation}, along with the signs of activity (spots) that we see manifested in \epic\ as quasi-periodic variability in the light curve. Then in Section~\ref{sec:models} we compare two sets of current stellar evolution models against our measurements of the physical properties (mass, radius, temperature), and illustrate the good agreement (with no adjustable parameters) of the same best-fit models with brightness measurements for other cluster members in the optical and near-infrared color-magnitude diagrams. We conclude in Section~\ref{sec:discussion} with our final remarks. | \label{sec:discussion} With our accurate determination of the masses, radii, and effective temperatures of its components, \epic\ joins the ranks of the eclipsing binaries with the best measured properties \citep[e.g.,][]{Torres:2010a}. The stars are slightly evolved from the zero-age main sequence, and both appear to be somewhat active, especially the primary. Despite the relatively short period of the system there is a hint of a very small orbital eccentricity. The theory of equilibrium tides \citep{Zahn:1977, Zahn:1989, Claret:1997} predicts a timescale for tidal forces to circularize the orbit of convective stars such as these of roughly 20~Gyr \citep{Hilditch:2001}, which would not be inconsistent with the observation. However, other empirical evidence suggests tidal forces are more efficient than predicted \citep[see, e.g.][]{Meibom:2005}, which would imply a shorter timescale. The comparison of our measurements against stellar evolution models in the mass-radius diagram yields fairly precise but model-dependent estimates of the age of $2.65 \pm 0.25$~Gyr for PARSEC and $2.48 \pm 0.30$~Gyr for MIST, with an additional uncertainty of 0.13~Gyr in each coming from the metallicity error. Predictions based on the same best-fit models also agree very well with our measured effective temperatures, which are less fundamental than the masses and radii. The temperatures, in turn, when combined with available standard photometry for the system, lead to an estimate of the extinction toward \epic\ that is similar to previous estimates but is determined in a completely different way. What distinguishes \epic\ from most other field eclipsing binaries with similarly accurate properties is its membership in a star cluster. This enables much more stringent tests of stellar evolution models, a better characterization of the parent population, and also valuable cross-validation of other techniques for dating stars based, e.g., on their oscillation frequencies or rotation periods. \epic\ belongs to Ruprecht~147 \citep{Curtis:2013}, a relatively nearby (300~pc) middle-aged open cluster observed by \ktwo, with a known metallicity determined spectroscopically from about half a dozen of its members. Using our estimate of the extinction, along with {\it Gaia}/DR2 parallaxes of cluster members, we find that the same model isochrones from PARSEC and MIST that reproduce the binary measurements in the mass-radius diagram also provide a reasonably good fit in the color-magnitude diagram of the cluster, particularly for the PARSEC model, both in the optical ($G$, $G_{\rm BP}$, and $G_{\rm RP}$ passbands from {\it Gaia}) and in the near infrared ($W1$, $J$, and $K_S$ passbands from WISE and 2MASS). The models therefore appear to be self-consistent in terms of their predictions of the internal structure, on the one hand, and the passband-specific fluxes, on the other, which depend on the adopted bolometric corrections, color/temperature transformations, and model atmospheres. We emphasize once again that the comparison in the CMD has no adjustable parameters: the age and extinction are determined separately based on the binary observations, the distance to all members has recently been supplied by the {\it Gaia\/} mission, and the metallicity is a known quantity. Few such tests have been available in the past because well-studied eclipsing binaries in clusters are rare. Remarkably, at least four additional relatively bright ($K\!p < 13$) detached eclipsing binaries that are amenable to study have been found in Ruprecht~147 and have high-quality space-based light curves, a situation that to our knowledge is unprecedented for an open cluster. \epic\ is the first of these eclipsing binaries to be analyzed. Spectroscopic observations of the others are well underway, and complete studies of them will be the subject of future publications by our team. Together, the components of these five binaries span a factor of two in mass ($\sim$0.7--1.4~$M_{\sun}$) reaching up to the turnoff region. Provided a similar precision as in \epic\ can be achieved for the masses and radii of the other four systems, a uniquely strong test of models in the mass-radius diagram will become possible, including perhaps a determination of the helium abundance in Ruprecht~147, given that the metallicity is known \citep[see, e.g.,][]{Brogaard:2011, Brogaard:2012}. We expect the age of the cluster could be determined considerably more precisely than we have here by using all five binaries simultaneously, and this could in turn serve as an important calibrator for other methods of estimating ages. For example, the \ktwo\ mission collected short-cadence photometric observations suitable for asteroseismic studies of about two dozen stars near the turnoff, as well as for a number of giants.\footnote{As reported in \url https://keplerscience.arc.nasa.gov/data/K2-programs/GO7035.txt~.} Seismic ages from these data, when they become available, could be compared against the binary age. Furthermore, \ktwo\ has provided a wealth of information for measuring rotation periods in Ruprecht~147. With a precise age determined from \epic\ and the other eclipsing binaries, the rotational sequence (mass-$P_{\rm rot}$ relation) for Ruprecht~147 could establish the cluster as a new reference point for gyrochronology \citep{Barnes:2007} at an age similar to that of NGC~6819 (2.5~Gyr), allowing for a valuable comparison between two old clusters \citep[see][]{Meibom:2015}. | 18 | 8 | 1808.07482 |
1808 | 1808.07161_arXiv.txt | {\it Fermi} has detected hard X-ray (HXR) and gamma-ray photons from three flares, which according to \stereo occurred in active regions behind the limb of the Sun as delineated by near Earth instruments. For two of these flares \r has provided HXR images with sources located just above the limb, presumably from the loop top (LT) region of a relatively large loop. {\it Fermi}-GBM has detected HXRs and gamma-rays, and RSTN has detected microwaves emissions with similar light curves. This paper presents a quantitative analysis of these multiwavelength observations assuming that HXRs and microwaves are produced by electrons accelerated at the LT source, with emphasize on the importance of the proper treatment of escape of the particles from the acceleration-source region and the trans-relativistic nature of the analysis. The observed spectra are used to determine the magnetic field and relativistic electron spectra. It is found that a simple power-law in momentum (with cut off above a few 100 MeV) agrees with all observations, but in energy space a broken power law spectrum (steepening at $\sim mc^2$) may be required. It is also shown that the production of the $>100$ MeV photons detected by {\it Fermi}-LAT at the LT source would require more energy compared to photospheric emission. These energies are smaller than that required for electrons, so that the possibility that all the emissions originate in the LT cannot be ruled out on energetic grounds. However, the differences in the light curves and emission centroids of HXRs and $>100$ MeV gamma-rays favor a different source for the latter. | \label{intro} {\it Fermi} Gamma Ray Observatory ({\it Fermi}; Atwood et al. 2009) observes the Sun once every other orbit. During the past solar active phase its Large Area Telescope (LAT) has detected $>100$ MeV photons from more than 40 solar flares. Few of these are detected only during the impulsive phase coincident with hard X-rays (HXRs) produced as nonthermal electron bremsstrahlung (NTB) and nuclear gamma-ray lines excited by high energy ions (mostly protons) (Ackermann et al.~2012). There is considerable evidence that the electrons are accelerated in a reconnection region near the looptop (LT) of the flaring loops (Masuda et al. 1994; Petrosian et al.~2002; Nitta et al.~2010; Krucker et al.~2010; Liu et al.~2013), and it is generally assumed that this is the site of acceleration of the impulsive phase protons (and ions) as well. But a majority of the LAT detected flares show only long duration emission (extending up to 10's of hours) usually rising after the impulsive phase (Ajello et al.~2014). Some stronger flares show both impulsive and gradual emission (Ackermann et al.~ 2014). Almost all LAT flares are associated with relatively fast ($>1000$ km/s) Coronal Mass Ejections (CMEs) and are often accompanied with gradual Solar Energetic Particle (SEP) events. This may indicate that the high energy particles responsible for the LAT gamma-rays are accelerated in the CME shock environment where the SEPs are produced. However, while SEPs are particles escaping the upstream region of the CME shock the gamma-ray producing particles, if originating at the CME, most likely come from the downstream region of the shock, with magnetic connection to the higher density solar atmosphere, which is the only place such high energy radiation can be produced. This scenario has received further support from \f-LAT detection of three flares which, as observed by {\it STEREO}, originate from active regions (ARs) behind the limb (BTL) of the Sun as delineated by near-Earth instruments. The analysis and some preliminary interpretation of the data from \f and other instruments on the BTL flares are presented in Pesce-Rollins et al.(2015) and Ackermann et al. (2017; Ack17). Our aim here is a more detailed modeling of the BTL flares with the particular focus on the determination of {\bf electron} spectra and energy contents required to produce the multiwavelength radiations seen in two of these flares. It should be noted that flares, such as these with occulted foot points, provide a clearer view of the coronal LT source, which may be the site of particle acceleration. Thus, the analysis presented below provides the most direct information on the acceleration process. There are several reports of observations of partially occulted flares in HXRs (see, e.g. Frost \& Dennis 1971; Krucker et al. 2007) and in gamma-ray emissions (Vestrand \& Forrest 1993; Barat et al. 1994; Vilmer et al. 1999). More recently Effenberger et al. (2017) have provided a complete list of \r observed partially occulted flares combining those from cycle 24 with the earlier list by Krucker \& Lin (2008) from cycle 23. Analysis similar to that presented here can be carried out for any of these flare with contemporaneous microwave coverage. The next section presents a summary of the relevant observational characteristics of these flares (all taken from Ack17). \S 3, provides a description of the main focus of this paper, which is to describe the emission processes and to determine the characteristics of the nonthermal {\bf electrons} required for their production. \S 4 contains a brief discussion of the possibility of LT origin of $>100$ MeV gamma-rays detected by the LAT. A summary and conclusions are presented in \S 5. | In this paper we present a detailed analysis of HXR and microwave spectra of two solar flares (Oct13 and Sep14) which, based on \stereo observations, originated 10 and 40 degrees BTL of the Sun, but were detected by {\it Fermi}, {\it RHESSI}, {\it SDO} \konus and ground based radio telescopes. The relevant observed characteristics are summarized in Tables 1 and 2. The 20-30 min HXR light curves observed by {\it RHESSI}, {\it Fermi}-GBM and \konus are almost identical and similar to the radio light curves for both flares. The {\it Fermi}-LAT light curves are somewhat delayed and last longer. \r images (up to 25-50 keV for Oct13 and 6-12 keV for Sep14) show sources (of size $\sim 50 \arcsec$) at the limb presumably the top of a relatively large flare loop peeking over the limb. The LAT localizations puts the centroid of gamma-ray emission $65 \arcsec$ and $275 \arcsec$ away from the \r source for Oct13 and Sep14 flares, respectively. Based on the similarity of light curves we assume a co-spatial emission of HXR and microwave emissions and determine accelerated electron characteristics over a broad range of energies from sub-relativistic regime (based on bremsstrahlung emission of low HXRs) to extreme relativistic regime (based on bremsstrahlung emission of gamma-rays and synchrotron emission of microwaves). In case of Sep14 flare the measured electron bremsstrahlung emission extends from 30 keV to $\sim 100$ MeV. This requires careful consideration of two important aspects. The first is the question of the time accelerated particles spend in the source region, which we call the escape time, and the second is that, because the observations span the trans-relativistic region, we should distinguish between spectra in momentum and energy space. Using empirically determined values and energy dependence of the escape time in 10-100 keV range by CP13 and P16, and their extension to relativistic energies based on theoretical considerations, we show that we are dealing with a thin target processes which then allows us to get the electron characteristics. Our results can be summarized as follows. \begin{enumerate} \item From modeling of the NTB emission of Oct13 flare we find that simple power law electron spectra in both momentum and energy space can reproduce the observed HXRs. For Sep14 flare a simple power law in momentum can describe the broad range of the observed HXRs more readily and with more reasonable values for the escape time index than simple power law in energy. From these fits we determine the spectral index, numbers and energy content of accelerated electrons. \item The radio spectra for both flares show a distinct optically thin emissions that peak around 1 GHz and a well defined turnover at lower frequencies indicating emergence of a optically thick spectrum. Self-absorbed synchrotron spectrum provides an adequate fit for the Sep14 flare, but for the Oct13 flare a self-absorbed synchrotron spectrum does not fit the observations in the range $0.5<\nu<5$ GHz. We show that a model whereby free-free absorption starts at about 7 GHz with self-absorption becoming dominant below 2 GHz provides an acceptable fit. These modelings allow us to determine both the spectrum and numbers of relativistic electrons and the magnetic field (that turn out to be lower than usual appropriate for the large height of the source). \item We then compare the two electron spectra obtained by these two methods. We show that for both flares extrapolation of spectra based on HXRs to the relativistic regime agree with those based radio data assuming a simple power law (with exponential cut off at several 100 MeV) in momentum but not in energy space. The latter require a broken power law with a break at $E<mc^2$. The numbers and energy content of these flares are in the right ball park and allow us to predict the FP emissions from AR located BTL. \item We also consider the possibility of thin target LT emission of the LAT gamma-rays and find that this requires 100 to 1000 time more energy of accelerated protons compared to thick target photospheric emission. However, even these energies are less than those of the electrons so that this scenario of high energy gamma-ray LT emission cannot be ruled out on energetic grounds. This is also true for production of these higher energy gamma-rays by GeV electrons at the photosphere. Nevertheless, because of the difficulty of acceleration of electrons to several GeV, pion decay scenario is favored, and the differences in the light curves and centroids of HXR and $>100 $ MeV emissions indicates a different acceleration site and mechanisms for (pion producing) protons than (HXR-microwave producing) electrons. \item The radiative signatures of occulted flares, such as those considered here, provide the most direct information on spectra and energy content of accelerated particles, and hence on the acceleration mechanism, uncontaminated by the stronger FP emission. For example, the differences between the required spectra in energy and momentum spaces can shed light on the details of the acceleration process. This important aspect of the problem will be dealt with in subsequent papers. \end{enumerate} {\bf Acknowledgements}: This work is supported by NASA LWS grant NNX13AF79G, H-SR grant NNX14AG03G and Fermi-GI grant NNX12AO78G. I would like to thank the {\it Fermi} colaboration, in particular the corresponding authors of Ack17 (A. Allafort, M. Pesce-Rollins, N. Omodei, F. Rubio and W. Liu) for help in preparation of this paper. I would also like to thank anonymous referees for many helpful comments. \vspace{1cm} | 18 | 8 | 1808.07161 |
1808 | 1808.02328_arXiv.txt | We investigate remnant neutron star masses (in particular, the minimum allowed mass) by performing advanced stellar evolution calculations and neutrino-radiation hydrodynamics simulations for core-collapse supernova explosions. We find that, based on standard astrophysical scenarios, low-mass carbon-oxygen cores can have sufficiently massive iron cores that eventually collapse, explode as supernovae, and give rise to remnant neutron stars that have a minimum mass of 1.17 M$_\odot$ --- compatible with the lowest mass of the neutron star precisely measured in a binary system of PSR J0453+1559. | The mass of neutron stars (NSs) is one of the most important observables to probe high-density nuclear physics. In particular, the maximum mass of NSs gives a stringent constraint on the nuclear physics above the saturation density. The current largest mass is $\approx$ 2 M$_\odot$ \citep{demo10,anto13}, which means that hypothetical nuclear equations of state having maximum NS mass smaller than 2 M$_\odot$ are excluded. Neutron star masses have a broad distribution. A precise measurement is possible for a binary system which contains at least one pulsar \citep{ozel16}. Recently, the first asymmetric system of double neutron stars, PSR J0453+1559, was discovered \citep{mart15}. The secondary NS's mass is much smaller than a canonical mass, that is, $1.174\pm0.004$ M$_\odot$. A corresponding baryonic mass of this NS is $\approx$ 1.28 M$_\odot$, which is remarkably smaller than a typical mass of Fe cores, i.e. $\sim$1.3--1.6 M$_\odot$ \citep[e.g.,][]{sukh18}. From this observation, a natural question arises: {\it Is it possible to form such a low-mass NS within the standard scenario of supernova (SN) explosion?} Although the explosion mechanism of core-collapse SNe is still unclear, there is the standard paradigm of neutrino-driven explosion \citep{beth85}, in which neutrinos produced in the vicinity of newly-born NSs heat up the post-shock material. Aided by the multi-dimensional hydrodynamic effects, e.g. convection and standing accretion shock instability, explosions driven by neutrino heating have been reported in the past decade \cite[see][for recent reviews and references therein]{burr13,jank16}.\footnote{ Note that the current simulations do not account for the observed explosion energy because they are not evolved for a sufficiently long time to see the convergence of the explosion energy. In addition, the nickel amount would be a more difficult to be explained by the current simulations because of the small growth rate of the explosion energy \citep{suwa17}.} In this work, we assume that SN explosion is driven by the neutrino heating mechanism. Producing a low-mass NS is also an issue for massive star evolution. Evolved stars with a carbon-oxygen (CO) core heavier than $\sim$1.37 M$_\odot$ have a possibility of SN explosion and NS formation \cite[e.g.][]{nomo87}. Single stars with a zero-age-main-sequence (ZAMS) mass of $\sim$8--12 M$_\odot$ have a path to electron-capture (EC) SNe \citep{nomo87,taka13} or core-collapse SNe from a low-mass Fe core \citep[e.g.][]{woos80,nomo84,nomo88,umed12,woos15}. Studies on a large number of progenitors suggest that these stars would produce low-mass NSs \citep{ugli12,ertl16,sukh16}. Note that fall-back accretion would increase NS mass for stars with a massive envelope, and thus, the formation mechanism of the low-mass NS is not trivial. In close-binary systems, an ultra-stripped SN is a possible path to produce a low-mass NS \citep[e.g.][]{taur13,taur15,suwa15,mori17,mull18}. Most of the H and He envelopes for these stars could be lost during their evolution by the close binary interactions. The explosion process of ultra-stripped SNe, especially the mass accretion history onto the proto-neutron star (PNS), would be considerably different from those of single stars. Thus, in this paper, we investigate NS formation in ultra-stripped SNe. In the following, we investigate the path to produce such a low-mass NS based on standard methods of stellar evolution and supernova explosion simulations. As in the previous works \citep{suwa15,yosh17}, which explored the evolution of massive stars whose envelope is supposed to be significantly stripped, we systematically study the dependence of the stellar evolution on the initial CO core mass in a parametric manner. By extending previous work toward lower CO core mass, we evaluate the minimum mass of an NS determined from stellar evolution and demonstrate that a low-mass NS like PSR J0453+1559 can be produced in the standard scenario of binary NS formation. The paper is organized as follows. Section \ref{sec:evolution} describes our stellar evolutionary calculations, in particular focusing on the consequent core masses. Section \ref{sec:coremass} gives estimates of Chandrasekhar mass which depends on the profiles of the electron fraction and entropy as well as the iron core mass. The numerical method of subsequent radiation hydrodynamics simulations and the results are presented in Section \ref{sec:explosion}. In Section \ref{sec:ECSN} we discuss differences of EC SNe from Fe-core forming core-collapse SNe in ultra-stripped SNe. We summarize our results in Section \ref{sec:summary}. | \label{sec:summary} In this paper, we investigated the minimum mass of NSs based on astrophysical scenarios, i.e., stellar collapse and supernova explosion. We calculated the evolution of CO cores with masses 1.35--1.45 M$_\odot$. The stars with CO-core masses higher than 1.36 M$_\odot$ form an Fe core after off-center Ne, O, and Si-burnings. We found that low-mass CO cores, which eventually form an Fe core and subsequently collapse, could result in an NS with mass $\sim 1.17$ M$_\odot$, which is comparable with the lowest NS mass precisely measured. \begin{figure} \centering \includegraphics[width=.45\textwidth]{Mco2Mns.eps} \caption{A schematic relation between NS mass and CO core mass. Two numbers in NS mass are baryonic mass ($M_{\rm NS,b}$; left) and gravitational mass ($M_{\rm NS,g}$; right), respectively. For $M_{\rm CO}\lesssim $1.37M$_\odot$, the electron-capture SN would be produced and the consequent NS mass would be $\sim$1.37M$_\odot$ and 1.24M$_\odot$ in baryonic mass and gravitational mass. For $M_{\rm CO}\gtrsim$ 1.37 M$_\odot$, the core-collapse SN would be produced and because of its lower-mass Fe core the consequent mass of the NS is smaller than models which produce EC SN.} \label{fig:mco2mns} \end{figure} According to our stellar evolution simulations, the minimum mass of CO cores that produce an Fe core is $\sim 1.37$M$_\odot$ and below this value a ONe core is formed, which would lead to EC SN instead. Because of its higher value of $Y_e$, an EC SN would produce more massive NSs than a core-collapse SN from an Fe core. Therefore, the minimum mass of NSs is expected to be determined by the core-collapse SN of a low-mass CO core (see Figure \ref{fig:mco2mns}). The range of the Fe core mass for single stars of initial mass 9--10 M$_\odot$ in \citet{woos15} is similar, or smaller, than ultra-stripped SN progenitors. Thus, the lowest-mass NSs may also be formed from the collapse of low-mass Fe cores. In the case of single stars, however, the reverse shock is produced when the shock wave arrives at the interface of the H-rich envelope. Then, the fall-back material accretes onto the collapsed core and increases NS mass. In the case of ultra-stripped SNe, on the other hand, the fall-back will be negligible because of very thin He layer and no H-rich envelope. Due to the same reason, the primary NS generated by the first SN explosion (not ultra-stripped SN) in binary systems would be incompatible to the light NS in PSR J0453+1559. This is because the secondary star would supply mass to the NS during its giant phase and increase the NS mass. Thus, low-mass Fe cores in the progenitors of ultra-stripped SNe, which are conjectured second explosions in close-binary systems, would be more favorable to form lowest-mass NSs. | 18 | 8 | 1808.02328 |
1808 | 1808.10017_arXiv.txt | Transmission spectroscopy provides a powerful probe of the atmospheric properties of transiting exoplanets. To date, studies of exoplanets in transit have focused on inferring their atmospheric properties such as chemical compositions, cloud/haze properties, and temperature structures. However, surface inhomogeneities in the host stars of exoplanets in the form of cool spots and hot faculae can in principle imprint signatures on the observed planetary transit spectrum. Here we present \textsc{Aura}, a new retrieval paradigm for inferring both planetary and stellar properties from a transmission spectrum. We apply our retrieval framework to a sample of hot giant exoplanets to determine the significance of stellar heterogeneity and clouds/hazes in their spectra. The retrieval analyses distinguish four groups of planets. First, the spectra of WASP-6b and WASP-39b are best characterised by imprints of stellar heterogeneity and hazes and/or clouds. HD 209458b and HAT-P-12b comprise the second group for which there is weak evidence for stellar heterogeneity and a high significance of hazes and/or clouds. The third group constitutes HAT-P-1b and WASP-31b and shows weak evidence against stellar heterogeneity but weak to substantial indications of clouds/hazes. The fourth group -- WASP-19b, WASP-17b, and WASP-12b -- is fit best by molecular and alkali absorbers with H$_2$ scattering without evidence for stellar heterogeneity and weak to no evidence for clouds/hazes. Our retrieval methodology paves the way to simultaneous information on the star and planet from higher resolution spectra using future facilities such as the James Webb Space Telescope and large ground-based facilities. | Transmission spectroscopy has been one of the most successful ways towards characterising exoplanetary atmospheres \citep[][]{burrows14,madhu16}. Studies of exoplanets in transit have been used to infer a wide variety of atmospheric properties including chemical compositions and abundances, clouds and hazes, and temperature profiles \citep[e.g.][]{sing16, kreidberg14a, kreidberg14b, knutson14a,knutson14b, madhu14, pont13, chen18, demory13, macdonald17, sedaghati18}. An essential assumption of most such studies is a homogeneous stellar photosphere characterized by one disk-integrated spectrum, and hence one stellar temperature and radius. However, stellar photospheres are not perfectly homogeneous. The stellar disk is generally comprised of differential areas each with a unique spectrum that can differ substantially from one representative average spectrum \citep{chapman87, shapiro14}. Active regions of the stellar surface, in the form of cool spots or hot faculae, are among the primary features of stars responsible for heterogeneity of the photosphere and variability of the stellar brightness in time. Such features on active stars may factor significantly and can induce modifications to an otherwise pristine planetary transmission spectrum and hence retrieved atmospheric properties \citep{rackham17, oshagh14, mccullough14, pont13,barstow15}. The slope of a transmission spectrum in the visible wavelength region is usually interpreted as hazes and/or clouds composed of small particles, and yet can also be caused by cool star spots which are unocculted by the transiting planet. For example, the steep optical slope of HD 189733b's transmission spectrum has been interpreted as opacity from haze particles in the planetary atmosphere \citep{pont13} as well as a possible signature of activity features in the photosphere of its variable host star, either unocculted star spots \citep{mccullough14} or occulted stellar plages \citep[i.e. bright chromospheric regions,][]{oshagh14}. On the other hand, hot faculae or plages, when not occulted by the transiting planet, decrease the observed optical transit spectrum due to an increasing stellar radius at shorter wavelengths. The optical transmission spectrum of GJ 1214b, which displays a significant decline towards shorter wavelengths and shallower transit depths than observed in the near-infrared \citep{kreidberg14a}, has recently suggested a heterogeneous stellar photosphere dominated by hot faculae \citep{rackham17}. Unocculted faculae have also been suggested to affect the transmission spectrum of GJ 1132b, which displays a significant decrease in transit depth at optical wavelengths \citep{dittmann17} and, like GJ 1214b, orbits a mid-M dwarf star \citep{berta-thompson15}. The suite of effects issued by active stars relates the importance for a framework capable of simultaneously dissecting properties of the planetary atmosphere and heterogeneous stellar photosphere from an observed spectrum. Here we present {\sc Aura}, a new retrieval framework that enables conjoint inference of exoplanetary and stellar properties. In addition to the usual properties explored for exoplanetary atmospheres (chemical abundances, clouds and hazes, and temperature structures), we incorporate a model that generally accounts for activity properties of the stellar photosphere \citep{rackham17}. The latter enables inferences on the fractional disk coverage of unocculted heterogeneity features, the average temperature thereof, and the temperature of the stellar photosphere. Our retrieval methodology ushers in the first instance of simultaneous retrieval of stellar and planetary properties constrained with present spectra, a consideration that will be integral to interpretations of future high fidelity observations from the James Webb Space Telescope (JWST) and large ground-based facilities. Our work is presented as follows. We detail the components of our retrieval framework in Section \ref{retrieval_model} and demonstrate the self-consistency of the retrieval methodology in Section \ref{retrieval_consistency}. We apply the new methodology to a sample of hot giant exoplanets in Section \ref{retrieval_application} to determine the significance of stellar heterogeneity and clouds/hazes in their observed transmission spectra. Limitations of our approach and necessary future considerations are presented in Section \ref{sec:limitations}. We discuss future observing prospects and summarise in Section \ref{discussion_conclusions}. | The four groups of planets identified in Table \ref{tab:condensed_expl} are approximately arranged by increasing chromospheric activity index (log $R'_{\mr{HK}}$), as seen by comparison of Tables \ref{tab:system_properties} and \ref{tab:condensed_expl}. The evident exceptions to this arrangement are WASP-19b and WASP-39b. Highly active stars (i.e., WASP-6) are found in the first group with evidence for stellar heterogeneity while low activity stars (i.e. WASP-12 and WASP-17) settle into the fourth group with substantial evidence against heterogeneity. The remaining planets (excepting WASP-19b and WASP-39b) are essentially ordered into the second and third groups, with weak evidence for or against heterogeneity. The observational consequences of this rough ordering of planets based on log $R'_{\mr{HK}}$ is noteworthy, for it suggests that the popular chromospheric activity indicator offers some predictive power as to whether an exoplanet transmission spectrum will be affected by stellar heterogeneity. The substantial and yet indefinite suggestions of stellar heterogeneity in the spectra of WASP-6b and WASP-39b offer definite observational strategies to increase the significance of inferred stellar heterogeneity effects and planetary properties with our retrieval methodology. While the effects of stellar heterogeneity can be significant at wavelengths as long as 2 $\mu$m (see Figure \ref{fig:het_effects}), they are most pronounced between 0.3 and 1 $\mu$m and thus high-impact observations should focus on this spectral range to probe the activity of stellar photospheres. The corollary is that infrared observations essentially ascertain atmospheric properties of the exoplanet alone. Joint studies of exoplanetary atmosphere compositions and heterogeneous stellar photospheres will benefit from precision observations in the Na and K absorption bands achievable with multiple HST orbits using STIS 430L and 750L grisms. Cool spots can masquerade as Na and K absorption features for a variety of stellar temperatures, metallicities, and gravities. In such cases and for observations of limited precisions (i.e. $\lesssim$100 ppm), degenerate explanations arise between Na and K abundances and stellar activity influence. Higher-precision observations in the Na and K bands can break this degeneracy, in turn enabling more precise estimates of heterogeneity covering fractions, alkali abundances, and haze properties. In addition, joint inferences will also benefit from high-resolution, broadband observations in the 0.3 to 1.0 $\mu$m window. This aspiration is already achievable through intelligent use of the VLT FORS2 spectrograph \citep[see for e.g.][]{sedaghati18}. Such finely-sampled observations can serve as a microscope into the heterogeneous coverage fraction since $\delta$ amplifies features of the heterogeneous spectrum (see Figure \ref{fig:het_effects}, upper panel). The positive correlation between the coverage fraction and heterogeneity temperature seen in Figure \ref{fig:simret} would also imply better constraints on the heterogeneity temperature with such VLT observations. The JWST can also contribute much to the joint retrieval of stellar and exoplanetary properties. The multi-wavelength capabilities facilitated by JWST can help separate stellar and planetary spectral imprints as the contrast between heterogeneous zones and the pristine stellar photosphere decreases with longer wavelengths as shown in Figures \ref{fig:stel_components}-\ref{fig:het_effects}. Thus, for example, spectra across 5 to 28 $\mu$m obtained with observing modes of the MIRI instrument will essentially reflect properties of the exoplanetary atmosphere alone. Simultaneous study of stellar and exoplanetary properties will also be possible. The time-series observing mode of the NIRCam instrument will permit stellar activity monitoring in the 0.6 to 5.0 $\mu$m range. The amplitude of time-series photometry can reveal lower limits on heterogeneous covering fractions \citep{rackham18} which can in turn be used as retrieval constraints using spectroscopy in the 0.6 to 5.0 $\mu$m range from NIRISS and NIRSpec modes. In summary, we have presented a new retrieval methodology, {\sc Aura}, that enables simultaneous inference of the properties of exoplanetary atmospheres and their host stars in the transiting configuration. The developed framework permits the inference of general in-homogeneous properties of the star including the stellar disk fraction covered by heterogeneity features, the average temperature of the heterogeneous fraction, and the temperature of the stellar photosphere. Jointly with the three stellar properties, the methodology permits the retrieval of a host of exoplanet atmosphere properties including the chemical compositions and abundances, attributes of clouds and hazes, and the temperature profiles throughout the atmosphere. Our methodology is the first joint retrieval framework suited for the extraction of properties of exoplanetary atmospheres and their host stars. As such it sets a precedence for more detailed joint analysis techniques of exoplanets and their stars in the future. We have applied our methodology to the transmission spectra of a sample of hot giant exoplanets to ascertain the influence of stellar heterogeneity and clouds/hazes in their spectra. The analysis distinguishes four groups of planets defined by the components needed to best explain their spectra. These four groups are illustrated in Table \ref{tab:condensed_expl}. In the first case, we find that the spectra of WASP-6b and WASP-39b are best fit with stellar heterogeneity as well as hazes and/or clouds. In the second case, there is marginal evidence for stellar heterogeneity effects and beyond substantial evidence for hazes and/or clouds in the spectra of HD 209458b and HAT-P-12b. The third group constitutes HAT-P-1b and WASP-31b and shows weak evidence against stellar heterogeneity but weak to substantial indications of clouds/hazes. In the fourth group three planets -- WASP-19b, WASP-17b, and WASP-12b -- are fit best by alkali and molecular absorbers with H$_2$ scattering without stellar heterogeneity and weak to no evidence of cloud/haze coverage. We emphasize that the suggestion against heterogeneity for WASP-19b is potentially due to the low data quality of the spectrum, and thus future studies of WASP-19b may suggest differently. We note that joint retrievals of the stellar photosphere and exoplanetary atmosphere rely on the assumptions that the model components are reasonably correct and reasonably complete. Presently, however, even state-of-the-art models for stellar heterogeneity are based on very limited knowledge; furthermore, the hypothesized haze particles can provide similar spectral signatures in the optical. Thus, future efforts must collect better data to break the degeneracy between stellar contamination and possible atmospheric aerosols. Ultimately, upcoming observatories will provide improved spectral resolutions and precisions useful for more definite and detailed joint analyses of transmission spectra. The simultaneous information on stellar and planetary properties facilitated through {\sc Aura} serves as a helpful tool in the analysis of present high-precision spectra and future high fidelity observations from the JWST and powerful ground-based facilities. | 18 | 8 | 1808.10017 |
1808 | 1808.00238.txt | {}{}{}{}{} % 5 {} token are mandatory \begin{abstract} %%%%%%%%%%%%%%%%%%%%% {{There is strong observational evidence that many active galactic nuclei ({\textbf{AGNs}}) harbour super-massive black holes ({\textbf{SMBHs}}), demonstrating multi-accretion episodes during their life-time.} In such \textbf{AGNs}, corotating and counterrotating tori, or strongly misaligned disks, as related to the central Kerr {\textbf{SMBH}} spin, can report traces of the \textbf{AGNs} evolution.} {Here we concentrate on aggregates of accretion disks structures, \textbf{r}inged \textbf{a}ccretion \textbf{d}isks ({\textbf{RADs}}) orbiting a central Kerr \textbf{SMBH}, assuming that each torus of the \textbf{RADs} is centered in the equatorial plane of the attractor, {tori are \emph{coplanar} and axi-symmetric}. Many of the \textbf{RAD} aspects are governed mostly by the spin of the Kerr geometry.} {We classify Kerr black holes (\textbf{BHs}) due to their dimensionless spin, according to possible combinations of corotating and counterrotating equilibrium or unstable (accreting) tori composing the \textbf{RADs}. The number of accreting tori in \textbf{RADs} cannot exceed $n=2$. We present list of $14$ characteristic values of the Kerr \textbf{BH} dimensionless spin $a$ governing the classification in whole the black hole range $0\leq a\leq M$, uniquely constrained by the {\textbf{RAD}} properties. } {The spin values are remarkably close providing an accurate characterization of the Kerr attractors based on the \textbf{RAD} properties. \textbf{RAD} dynamics is richer in the spacetimes of high spin values. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% One of the critical predictions states that a \textbf{RAD} tori couple formed by an outer accreting corotating and an inner accreting counterrotating torus is expected to be observed only around slowly spinning ({$a<0.46M$}) \textbf{BHs}. The analysis strongly binds the fluid and \textbf{BH} characteristics providing indications on the situations where to search for \textbf{RADs} observational evidences. Obscuring and screening tori, possibly evident as traces in X-ray spectrum emission, are strongly constrained, eventually ruling out many assumptions used in the current investigations of the screening effects. % } %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% { We expect relevance of our classification of Kerr spacetimes in relation to astrophysical phenomena arising in different stages of \textbf{AGNs} life that could be observed by the planned X-ray satellite observatory \textbf{ATHENA} (\textbf{A}dvanced \textbf{T}elescope for \textbf{H}igh \textbf{EN}ergy \textbf{A}strophysics).} | %%%%%%%%%%%%%%%f%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%\input{intro} Black hole (\textbf{BH}) physics and the \textbf{BH} accretion disk investigation is developing in the last years. The launch of new satellite observatories in the near future allows an unprecedented close look at situations and contexts which were inconceivable only a few years ago. Observable features on the accreting disks, focusing on morphology of accretion processes and associated jet emissions provide increasingly more detailed and focused pictures of these objects. The sensational opening of a new observational era represented by gravitational waves (\textbf{GWs}) detection allows us to focus on questions of broader interest involving more deeply the \textbf{BH} physics. On the other hand, the theoretical modeling seems to resort deeply to these improvements. Concerning the \textbf{BH} accretion disks processes, there is great expectation towards the X-ray emission sector, with several missions as \textbf{XMM-Newton} (X-ray Multi-Mirror Mission)\footnote{{\textbf{http://sci.esa.int/science-e/www/area/index.cfm?fareaid=23}}}, \textbf{RXTE} (Rossi X-ray Timing Explorer)\footnote{{\textbf{http://heasarc.gsfc.nasa.gov/docs/xte/xtegof.html}}} or \textbf{ATHENA}\footnote{\textbf{http://the-athena-x-ray-observatory.eu/}}. Recent studies point out also an interesting possible connection between accretion processes and \textbf{GWs} \citep{PRL}. In this scenario, however, many questions remain still unresolved, leaving them as intriguing problems of observational astrophysics, appearing to demand even a greater effort from the point of view of the model development, as the still missing solution of \textbf{g}amma \textbf{r}ay \textbf{b}ursts (\textbf{GRBs}) origin, the jet launch, the \textbf{q}uasi-\textbf{p}eriodic \textbf{o}scillations (QPOs), the formation of \textbf{SMBHs} in \textbf{AGNs}. In general, the theoretical investigation is increasingly oriented towards the attempt to find a correlation between different phenomena and a broader embedding environment, {creating a general framework of analysis envisaging the \textbf{\textbf{BH}}-disk system as a whole.} In this sense we may talk about an {Environmental} {Astrophysics}. Evidences of this fact are the debates on the jet-accretion correlation, the \textbf{BH} accretion rate-disk luminosity issue, the \textbf{BH} growth--accretion disk and \textbf{BH}--spin shift--accretion disk correlation and \textbf{BH} populations and galaxy age correlation--see for example \citet{Hamb1,Sadowski:2015jaa,Madau(1988),Ricci:2017wmr,Narayan:2013gca,Mewes:2015gma,Morningstar:2014hea, Yu:2015lqj,Volonteri:2002vz,Yoshida,Regan:2017vre,Yang:2017slb,Xie:2017jbz}. In this regard, a crucial aspect to establish correlation between \textbf{SMBHs} and their environment is first the recognition of the \textbf{BH} attractor. Although the unambiguous \textbf{SMBH} identification reduces to assign the \textbf{BH} spin $a$ and mass $M$ parameters or even, for many purposes, { only the \textbf{BH}} spin-mass ratio $a/M$, this task is still controversial and debated issue of the \textbf{BH} astrophysics. We shall see in fact that many aspects of the \textbf{BH} accretion disk physics depend only on the ratio $a/M$. The issue to identify a rotating \textbf{BH} by determining its intrinsic rotation, or spin-mass ratios, is a rather difficult issue, a complex task which is challenged by different observational and theoretical approaches. All these methods are continuously debated and confronted--see for example \citet{Capellupo:2017qpt,McClintock:2006xd,Daly:2008zk}. It should be noted then that the evaluation of the \textbf{SMBHs} spin is strictly correlated with the ``mass-problem'': the assessment of the precise value of the spin parameter of the \textbf{BH} is connected with the evaluation of the main features of the \textbf{BH} accretion disk system, as the \textbf{BH} accretion rate or the location of the inner edge of the accretion disk. The \textbf{GWs} detection from coalescence of \textbf{BHs} in a binary system may serve, in future, as a further possible method to fix a \textbf{BH} spin parameter \citep{Farr:2017uvj,vanPutten:2016wpa,vanPutten:2012vr,vanPutten:2015eda}. However, nowadays this task is often approached in \textbf{BH}-accretion disk framework, by considering the evaluation of the mass accretion rate (connected with the disk luminosity) or, for example, the location of the inner edge of an accreting disk. Nevertheless all these aspects are certainly not settled in one univocal picture; for example, even the definition of the inner edge of an accreting disk is very controversial -see for example \citet{Krolik:2002ae,BMP98,2010A&A...521A..15A,Agol:1999dn,Paczynski:2000tz,open,long}. {Therefore, GW observations have the ability to provide a ``disk''-independent way to trace back the (dimensionless) \textbf{BH} spin and more generally an evaluation of the mass and spin parameters of the black holes -- for GW methods for the evaluation of the \textbf{BH} spin see for example \cite{Farr:2017uvj,vanPutten:2015eda,vanPutten:2016wpa,2016PhRvD..93h4042P, 2016ApJ...826...91A}. } In this paper, we face the problem of the black hole identification, proposing an approach which we believe can be promising especially for \textbf{SMBHs} in \textbf{AGNs}. Our investigation is essentially centered on the exploitation of a special connection between \textbf{SMBHs} and their accretion tori; we show that the dimensionless spin of a central \textbf{BH} ($a/M$) and the morphological and equilibrium properties of its accreting disks, are strongly related. Recently, various analyses have shown, in the general relativistic regime, that completely axis-symmetric and coplanar configurations are strongly restricted with regards to their formation, kinematic characteristics (as range of variation of angular momentum) or the emergence of instability. Their existence is constrained according to different evolutionary phases of the individual configurations and, more importantly, by the properties of the central Kerr attractor \citep{pugtot,ringed,open,dsystem,long}. More generally, there is a strong relation between the galaxy dynamics and its super-massive guest, specially in the accretion processes. It is expected that such \textbf{SMBHs} in \textbf{AGNs} are characterized by a series of multi-accreting episodes during their life-time as a consequence of interaction with the galactic environment, made up by stars and dust, being influenced by the galaxy dynamics. Further example of the complex and rich black hole-active galaxy interaction is known as \emph{feedback \textbf{AGN}}-- \citep{Ricci:2017abh,Ricci,Ricci:2017wmr,Yoshida,Komossa:2015qya}. These activities may leave traces in the form of matter remnants orbiting the central attractors. Thus, chaotical, discontinuous accretion episodes can produce sequences of orbiting toroidal structures with strongly differing features as, for example, different rotation orientations with respect to the central Kerr \textbf{BH} where corotating and counterrotating accretion stages can be mixed \citep{Dyda:2014pia,Aligetal(2013),Carmona-Loaiza:2015fqa,Lovelace:1996kx,Romanova}. Strongly misaligned disks may appear with respect to the central \textbf{SMBH} spin \citep{Nixon:2013qfa,2015MNRAS.449.1251D,Bonnerot:2015ara,Aly:2015vqa}. However, in this work classes of rotating \textbf{BHs}, identified by their dimensionless spin, are associated to particular features of the \textbf{BH} orbiting accreting tori, namely aggregates of toroidal axis-symmetric accretion disks, also known as Ringed Accretion Disks (\textbf{RADs}), centered on the equatorial plane of a Kerr \textbf{SMBH}. The \textbf{RAD} aggregate is composed by both corotating and counterrotating tori, the limiting case of single accretion torus orbiting the \textbf{SMBH} is also addressed as a special case of the \textbf{RAD}. Each \textbf{RAD} toroidal component is modeled by a perfect fluid with barotropic equation of state and constant specific angular momentum distribution \citep{Pu:mnrasPD,pugtot,abrafra,Pugliese:2013hp}. \textbf{RAD} models follow the possibility that several accretion tori can be formed around very compact objects as \textbf{SMBHs} ($10^6-10^9 M_{\odot}$, $M_{\odot}$ being solar masses) in \textbf{AGNs}. \textbf{RAD} may be also originated after different accretion phases in some binary systems or \textbf{BH} kick-out, or by local clouds accretion. {We mention also \cite{Bonnell} for an analysis of the massive cloud spiraling into the \textbf{SMBH} Galaxy.} {More generally, multiple structures orbiting around \textbf{SMBHs} can be created from several different processes involving the interaction between the \textbf{BH} attractors and their environment. An original Keplerian disk can split into two (or more) components (toroids) for some destructive effects where for example the \emph{self-gravity } of the disk becomes relevant. The impact of the disk self-gravity in different aspects of the \textbf{BH} accretion is discussed in Sec.\il\ref{Sec:gothc}. Formation of more tori is also one of the possible endings of a misaligned disk, for example in a binary system, where the torque and warping is relevant to induce a disk fragmentation. In all these cases, the rotational law of newly formed tori can be also very different. In fact, multiple systems can originate in several periods of the accretion life from different material embeddings.} {In an early phases of their evolution, accretion disks can be misaligned with respect to the equatorial plane of the Kerr attractor and in many cases such disks are expected to be ``warped'' and ``twisted'' accretion disks. Although the misaligned or warped case is not covered here, we also discuss the occurrence of this possibility within the \textbf{RAD} frame in Sec.\il\ref{Sec:gothc}, where we also address possible instability processes when \textbf{RAD} is extended to consider aggregates with the contribution of the magnetic field. However, misaligned disks, for specific values of the characteristic parameters, will eventually end in a steady state with an inner aligned disk. It has been shown that counterrotating tori can derive also from highly misaligned disks after galaxy mergers, with galactic planes strongly inclined. } {Furthermore, in a warped disk scenario the analysis of inner region of accretion disk connected with the jet emission is also considered for the assessment of the central \textbf{BH} spin, assuming that the jet directions are indicative of the direction of the \textbf{BH} spin. We consider this briefly in Sec.\il\ref{Sec:gothc}. In the case of misaligned disks, these studies focus on radio--jet direction in \textbf{AGN} along orbital plane direction. More generally, \textbf{BH} and jet emission connection (for example radio and X-ray emission) are considered in galactic embedding to test correlation between the galactic host and jets; the presence of a warped disk can also explain the jet emission orientation with respect to the galactic plane showing also a strong misalignment. Without going into details of this aspect of the accreting process, going beyond the scope of the present work, we mention in particular the case of galaxy \textbf{NGC4258} %, connecting with \citep{Moran:2008dv, Humphreys:2007ir, Doeleman:2008qh, Qin:2008uu, Rodriguez:2008bi, Kondratko:2008gr}.} {Jet emission, in fact, constitutes the third ingredient in the unified \textbf{BH}--accretion disk framework. Almost any \textbf{BH} is associated to a jet emission. How exactly the jet emission and morphology (collimation along the axis, rapidity of emission launch, chemical composition) can be precisely linked or induced by the mechanisms of accretion remains to be clarified. Jets are supposed to be connected with the dynamics of the inner region of the disk in accretion. The role of jet in the \textbf{BH}-accretion disk systems and accretion physics is multiple, altering the energetic of the accretion processes with the extraction of the rotational energy of the central \textbf{BH} and the rotational energy of the disk, and changing the accretion disk inner edge. The inner--edge-jet correlation has to be then framed in the case of the misaligned disks, where the location of the edge is related to the strength of the jet-- \cite{NL2009,Fender:2004aw,Fender:2009re,Soleri:2010fz,Tetarenko18,AGNnoBH,FangMurase2018,1998NewAR..42..593F,1998MNRAS.300..573F,1999MNRAS.304..865F,Fender(2001)}. For an analysis of the energetic X-ray transient with associated relativistic jets, updated investigations on jet emission detection see \cite{NatureMa,Maraschi:2002pp,Chen:2015cga,Yu:2015lqj,Zhang:2015eka,Sbarrato:2014uxa,MSBNNW2009, Ghisellini:2014pwa,BMP98}. For jet-accretion disk correlation in \textbf{AGN} see for example \cite{liska,Caproni:2017nsh,Inoue:2017bgt,Gandhi:2017dix,Duran:2016wdi,Vedantham:2017kyb,Bogdan:2017pgi,Banados:2017unc,DAmmando:2017ufp} and also \cite{NL2009,Fender:2004aw,Fender:2009re,Soleri:2010fz,Tetarenko18,AGNnoBH,FangMurase2018}. } {Jet production, as an important additional aspect of the physics of accretion disk, fits into the \textbf{RAD} context in many ways. Firstly, more points of accretion can be present in the \textbf{RAD} \emph{inside} the ringed structure. It has been shown in \cite{ringed} that a \textbf{RAD}, as an aggregated body of orbiting tori, can be considered as a single disk orbiting around a central Kerr attractor on its equatorial plane, with axi-symmetric but knobby surface and diversified specific angular momentum distribution, due to the different contributions of each torus of the \textbf{RAD} agglomerate that can be either corotating or counterrotating with respect to the central attractor. Only for the \textbf{RAD} systems satisfying certain constraints on the \textbf{BH} spin and on the specific angular momentum of the tori, a double accretion phase can occur. Such a situation implies the concomitant presence of two coplanar accreting tori of the \textbf{RAD} in the same period of the \textbf{BH} life. In the frame of accretion disk--jet correlation, this would imply the presence of a shell of double jets, one from an outer counterrotating torus and one associated to the inner corotating accreting torus. Moreover, the geometrically thick tori considered as \textbf{RAD} components are known to be associated to several species of open surfaces (proto-jets) related to emission of matter funnels in jets--\cite{abrafra}. These configurations have been discussed in literature in several contexts. The \textbf{RAD} model inherits this characteristic of the accretion torus. The proto-jets, as related to the critical points of the hydrostatic pressure in the force balance of the \textbf{RAD} tori, can give raise to complicated sets of jets funnels either from corotating or counterrotating fluids in more points of the tori agglomeration. We do not consider directly in this investigation the open surfaces, but they will be considered for the classification of the attractors. A more focused analysis on proto-jets in the \textbf{RAD} framework can be found in \cite{open,long,app}. } The plan of this article is as follows. Section \il\ref{Sec:Taa-DISK} introduces the model of \textbf{RAD} orbiting a central Kerr attractor: in Sec.\il\ref{Sec:prop} we examine main properties of the Kerr \textbf{BH} exact solution. \textbf{RAD} model is then developed in Sec.\il\ref{Eq:tori} where the toroidal configurations are discussed, and main properties and characteristics of the tori and the macro-structure are presented. The \textbf{RAD} is a relatively new model, this analysis has therefore required the introduction of new concepts adapted to the model. For this purpose it is convenient to introduce a reference Section\il\ref{Sec:notation-sec} where we give relevant model details grouping together the main definitions used over the course of this work. We made use also of Table\il\ref{Table:pol-cy-multi} and Table\il\ref{Table:def-flour} where the symbols and relevant notation used throughout this article are listed. Section\il\ref{Sec:RADs-sis} encloses the main results of this work. Tables\il\ref{Table:nature-Att} shows the major classes of attractors considered in this section and, in a compact form, main properties the associated \textbf{RADs}, as investigated in Sec.\il\ref{Sec:notation-sec}. {In Sec.\il(\ref{Sec:nw}) we concentrate on the double accretion occurring in \textbf{RAD} couples}. In Sec.\il\ref{Sec:gothc} general discussion on the \textbf{RAD} instabilities and generalizations of the \textbf{RAD} model is presented. We also discuss the possible significance of the tori self-gravity, tori misalignment, viscosity and magnetic field in the \textbf{RAD}. Concluding remarks are in Sec.\il\ref{Sec:Open-Concl}, where we report a brief summary of the results of our analysis, followed by considerations on the impact of the \textbf{RAD} hypothesis in the \textbf{AGN} environments and on some aspects of the phenomenology connected with the \textbf{RAD}. Finally, Appendix follows, where further details on \textbf{BH} spin properties are provided. % | \label{Sec:Open-Concl} % We investigated agglomerations of accreting coplanar tori (\textbf{RADs}) orbiting in the equatorial plane of a Kerr \textbf{BH} attractor as tracers of the Kerr \textbf{SMBHs}. The limit of the \textbf{RAD} constituted by one torus in accretion is also considered. This analysis eventually proves the existence of a strict correlation between the \textbf{BH} spin-mass ratio $a/M$, and the accretion tori featured by the \textbf{RADs}. Our findings resulted in the complete description of the \textbf{RAD} systems characterizing the central Kerr \textbf{BHs} in classes which are uniquely identifiable through the properties of the rotating tori of the agglomerate. Overview of the major features of the classes of spinning attractors has been discussed, according to the properties of the orbiting toroidal structures. Table\il\ref{Table:nature-Att} collects the main classes characteristics. More specifically, we developed a general classification of attractors and orbiting tori gathering together information on the Kerr \textbf{BHs} spin and their orbiting accreting tori in the \textbf{RADs} framework introduced in \citet{pugtot,ringed,open,dsystem,long}. As consequence of this analysis, different issues related to the \textbf{SMBH-RAD} systems were addressed such as the identification of observable features of the toroidal agglomerates the indication of the associated \textbf{BH} classes according to their dimensionless spin. These classes associate a \textbf{RAD} to its central attractor, therefore they provide indication on the \textbf{BH} environments where it would be possible to observe a \textbf{RAD}. Eventually we addressed more explicitly the question whether there is a way to unambiguously identify the \textbf{RAD} aggregates through the specification of the \textbf{RADs} order, angular momentum distributions, the relative rotation of the toroids, toroid rotations with respect to the attractor, and internal dynamics of the \textbf{SMBHs} (the spin-mass ratio). This investigation serves to different purposes: to envisage an accretion disk--\textbf{SMBH} correlation in the \textbf{RADs} scenario, or in the limiting case of one orbiting torus, and to provide an indication of the possible attractor-disk candidates to look for observational evidences of \textbf{RADs}. This work supplies the most comprehensive reference template of \textbf{BH}-accretion tori correlation, which can be also used an a strictly constrained set of the initial configurations for the development of a fully general relativistic dynamical \textbf{GRMHD} simulation for the \textbf{SMBH}-\textbf{RAD} system--see for example \citet{Luci,abrafra}. Our analysis is conducted by numerical integration of the hydrodynamic equations for multiple tori with fixed boundary conditions for each configuration of the set--see Figs\il\ref{Fig:but-s-sou}. Parallel to this analysis we carefully explored the ranges of fluid specific angular momentum and $K$-parameters, drawing then the constraints for the radii $(r_{cent},r_{\times})$ and $(r_{in},r_{out},r_{coll})$--see also \citet{open,dsystem,app}. It is clear that the construction of the aggregates relies essentially on the boundary conditions imposed on the function (\ref{Eq:def-partialeK}) (the Heaviside functions). We used the studies of \citet{open,long}, which fix the location of the accretion torus edges in the spacetime regions confined by marginally bound, marginally stable and marginally circular (photon) orbits. This setup turns to be very important for the analysis of the \textbf{RAD} oscillations emerging as perturbations of tori and the boundary conditions \citep{ringed}. The \textbf{RAD} model encloses a huge amount of possibilities to be investigated, having a large number of cases even in the simple three parameters model considered here (the specific angular momentum $\ell$, the $K$ parameter and the attractor spin-mass ratio). Interestingly, this investigation enlightened also the importance of the dimensionless quantities $\ell/a$, $r_{cent}/a$ and $r_{in}/a$, while a deeper analysis of the \textbf{BHs}-accretion tori connection in this special parametrization is reserved for a planned future work. We consider these configurations using a seed to generate a general sequence of the \textbf{RAD} tori, starting from one of the four germs $\pp_{\pm}> \pp_{\pm}$ (the $\ell$corotating couples) or $\pp_{\pm}> \pp_{\mp}$ (the $\ell$counterrotating couples) respectively--Figs\il\ref{Figs:GROUND-scheme}. Focusing here on the central objects around which these configurations may orbit, and considering different properties of the configurations, we drawn a detailed classification of the attractors in \textbf{$17$} general classes singled by the their \textbf{BHs} spin-mass ratios in the entire range $0\leq a\leq M$, including the Schwarzschild static solution and the extreme Kerr solution. As expected, this classification tends to strongly differentiate fast spinning attractors and slow spinning ones. Major differences are highlighted for two classes of attractors, with spin $0< a\leq 0.47M$ and $a>0.47 M$ respectively. The static case, represented by the Schwarzschild solution, was considered separately. %The A further relevant aspect of this analysis is that the \textbf{BH} classes often intersect--see Fig.\il\ref{Fig:Relevany}. Importantly we have taken into account the \textbf{RAD} stability properties in the construction of the classes of attractors. The main aggregate instabilities are mainly driven by two families of processes: {\textbf{(1)}} Collisions between accreting or non-accreting tori, and \textbf{(2)} \textbf{RAD} instability following the accretion phase of one torus of the aggregate \citep{ringed,dsystem}. A double accretion can be observed \emph{only} in a couple $\cc_{\times}^-<\cc_{\times}^+$ (es Fig.\il\ref{Fig:but-s-sou}), around all Kerr \textbf{BHs} ($a\neq0$). Moreover, in the \textbf{RAD} scenario, the maximum number of accreting tori orbiting around one central Kerr \textbf{SMBH} is $n=2$. This opens up important potential observations encouraging also a review of the current interpretation of the accretion data by considering a \textbf{RAD} framework. In fact, the ringed structure can be effectively disguised as a geometrically thin, axi-symmetric disk centered on the equatorial plane of the Kerr \textbf{SMBH} with interrupted phases of super-Eddington accretions and a very rich inner dynamics with jet emission featuring also inter-tori, proto-shell, jet emission. The presence of an inner torus can also enter as a new unexpected ingredient in the accretion-jet puzzle--see also \citet{Hamb1}. Therefore, in Table\il\ref{Table:nature-Att} we also consider the possibility of launch of proto-jet configurations in the \textbf{RADs}. This situation however, has been deeply analyzed in \citet{open,long}. We have shown an important restriction on any screening effect from an inner torus of an aggregate. Screened X-ray emission by some ``bubbles'' of material are currently studied in many processes-\citep{Gilli:2006zi,Masini:2016yhl,DeGraf:2014hna,Storchi-Bergmann, Ricci:2017wmr,Ricci,Marchesi,Marchesi:2017did}. Then, particularly some \textbf{AGNs} have been proved to be obscured according to the X-ray spectral emission \citep{Marchesi,Marchesi:2017did}; several analyses also suggest that obscuration in optical and X-ray emission profile may be due to different phenomena, caused by dust materials surrounding the inner part of the galactic nuclei. So far these (free) dust materials were almost always supposed to be randomly distributed around the central \textbf{BH} and, depending on the gas density, the light emitted during the growth could be absorbed in the optic as in the X-ray electromagnetic band, distinguishing \textbf{AGN} as obscured, not obscured, or much obscured (or \emph{Compton thick}). We here explicitly claim for an analysis of this obscuration assuming a \textbf{RAD} scenario, therefore considering the inner corotating, accreting or quiescent torus in accordance with the analysis in Sec.\il\ref{Sec:RADs-sis}, and with the constraints imposed by the spin of the central \textbf{SMBHs}. Such \textbf{RADs} are very much constrained so that they cannot be randomly distributed neither assumed to exist independently on the evolution of the \textbf{SMBH} itself--see also \citet{dsystem,long}. We believe this change of paradigm may have a huge impact on the current scheme adopted to explain the obscuration. To be more precise, screening tori may only exist as corotating fluids in particular \textbf{BH} classes and under specific restrictions on the specific angular momentum. These tori can be accreting onto the central \textbf{BH} or quiescent-- see also \citet{dsystem,long}. A ``screening'' \emph{corotating}, non-accreting, torus between the two accreting tori can be observed only as $\,\cc_{\times}^-<\cc^-<...<\cc_{\times}^+$ for any Kerr attractor having spin $a\neq0$. These special tori are expected to be relatively small compared to the outer tori of the \textbf{RAD} agglomerate. A further case is where the inner corotating or counterrotating (for \textbf{BHs} with $a<0.46M$) accreting torus of the \textbf{RAD} is ``obscured'' by an \emph{outer} screening and quiescent torus. As seen in Table\il\ref{Table:nature-Att}, a counterrotating accreting torus with an outer corotating torus towards the accretion (i.e. a $\cc^-_1$ torus having specific angular momentum $\ell\in \mathbf{L1}^-$), can be observed only as a $\cc_{\times}^+<\cc_1^- $ aggregate, and orbiting around slow spinning \textbf{SMBHs} with spin $a<0.46M$. In the \textbf{BH} classification, we have also included very fast spinning \textbf{SMBHs} with $a>a_1$, which are characterized by the possibility that corotating (screening, accreting) tori can be orbiting in the ergoregion. This aspect, also mentioned in \citet{pugtot,ergon,ringed,dsystem,open}, is an interesting consequence of the frame dragging. {More generally, we expect \textbf{RAD} to be a relevant feature of the faster-spinning \textbf{SMBHs}. Specifically the $\ell$corotating \textbf{RADs} are favored features of slow attractors $(a\lesssim0.45M)$, while around faster spinning attractors, $\ell$counterrotating \textbf{RADs} would be most expected. Generally, in the spacetimes of the slower rotating \textbf{BHs}, tori collision appear more probably. This episode may be followed by the collisional energy release and, eventually, tori merging and accreting onto the central \textbf{BH}--\cite{app}. For the slow spinning \textbf{BHs}, with dimensionless spins $a/M$ close to the lower limit of the static attractors, any \textbf{RAD} tori couple has to be considered as an $\ell$corotating couple, independently of the relative rotation of the tori, i.e., although the tori of the aggregate can have alternate spin, eg. $(0,+,-)$ or $(0,-,+)$, according to the notation introduced in Fig.\il\ref{Figs:GROUND-scheme}, their properties considered for \textbf{RAD} structure, are entirely analogue to the $\ell$corotating case in a Kerr spacetime. Moreover, in the Schwarzschild spacetime, double accretion or obscuration is \emph{not} possible. On the other hand, very fast-spinning \textbf{SMBHs} $(a\approx M)$, would favor the formation of $\ell$counterrotating tori which are also largely separated in the \textbf{RAD}. This holds particularly for the couples $\pp^-<\pp^+$, where a double accretion phase occurs. Instead, the seed having an inner counterrotating torus in accretion with an outer corotating torus is favored only in the early phases of formation of the outer torus, for fast-spinning \textbf{SMBH}, while in Schwarzschild spacetime such a couple can be always observed--\citep{dsystem}. Kerr \textbf{SMBH} attractors are significant environments for the \textbf{RAD} observation. In fact, \textbf{SMBHs} determine certainly strong curvature effects having ultimately a major influence in the determination of the constraints in (\ref{Eq:def-partialeK}). Then, as shown in \citet{dsystem} and then discussed in \citet{long}, very massive \textbf{BHs} would prevent the emergence of tori collision. Note also that the space-scales are here in units of the \textbf{SMBH} mass; for example the maximum spacing $\bar{\lambda}$ considered for two accreting tori is $\bar{\lambda}\approx8 M$ in the case of nearly extreme \textbf{BHs}. To realize the significance of this we should note that the spacing parameter $\bar{\lambda}$ between two tori regulates in fact the possibility of tori collision. } { The unstable phases of the aggregate, on one hand, could undermine the survival of the ringed structure, eventually leading to the formation of a single disk. On the other hand, the \textbf{RAD} unstable phases constitute environment for set of interesting phenomena-- see \cite{dsystem}. It is possible that a \textbf{RAD} represents a final steady state of life of an attractor-disk system originating as extended disk misaligned with the \textbf{BH} axis; this hypothesis stands as particularly promising for the very massive \textbf{BH} attractors in \textbf{AGNs}, as demonstrated in Sec.\il\ref{Sec:gothc}--see also \cite{Nelson:2000mw,1996MNRAS.282..291S,2005MNRAS.363...49K,liska,2018arXiv180110180S, 2015MNRAS.449.1251D,2006MNRAS.368.1196L,1999ApJ...525..909A,2017ApJ...845...10F, 2009MNRAS.399.2249P,2009MNRAS.400..383M}. } \subsection{Relevance of the \textbf{RADs } in the \textbf{AGN} environments } It is clear that a \textbf{SMBH} is rarely an isolated \textbf{BH}, but rather a living object in the host galaxy cores interacting with its environment and subjected to several evolutionary phases, where its mass and spin will often change during different stages of the \textbf{BH} life. The \textbf{SMBHs} intercept the stellar population and dust of their galactic host and vice versa \textbf{SMBHs} remix, the matter and radiation content of the galactic environment. Suffering from galaxy collisions, the interaction between the \textbf{BH} and the galaxy environment ends in changes of the \textbf{BH} spacetime which is initially considered as ``frozen background''. A non-isolated \textbf{BH} background can change following a spin-down or a spin-up process \citep{Abra83,Abramowicz:1997sg,Rez-Zan-Fon:2003:ASTRA:,Font:2002bi,Hamersky:2013cza,Adamek:2013dza,Lot2013}. This clearly should lead to a \textbf{BH} shift from one class of our classification to another, and to a change of the equilibrium conditions of the \textbf{RADs}. Toroidal structures might be formed as remnants of several accretion regimes occurred in various phases of the \textbf{BH} life { \citep{Aligetal(2013),Lovelace:1996kx,Romanova,Volonteri:2002vz,Carmona-Loaiza:2015fqa,Dyda:2014pia}} Therefore, the analysis considered here can be significant for the resolution of the so called mass problem for \textbf{SMBHs} in \textbf{AGNs}. We can consider accretion in long and continuous accretion episodes, arising due to merging involving a large spin-shift or, \emph{vice versa}, sequences of small and random accretion episodes, being advocated for the formation of the \textbf{SMBHs}. The \textbf{RAD} internal dynamics has several inter-disk effects including double accretion processes, screening effects, and tori collisions. These situations clearly represent mechanisms for the mass growth of the \textbf{SMBHs}. It is clear that many aspects of the physics of \textbf{SMBHs} and their host galaxies would be altered by the relevance of the \textbf{RAD} model in support of the hypothesis of a more complex \textbf{BH}-accretion disk system, than is commonly considered. In Sec.\il\ref{Sec:gothc} we also examined the impact of the \textbf{RAD} scenario, regarding, more widely, future generalizations of the tori aggregate model to consider misaligned tori and, more generally, diverse accretion disk models. It should be also stressed that, although the current theoretical analysis conceives on a large variety of accretion models, with a diversified parametrization and disk shape (depending on disk optically depth, the geometric thickness, luminosity...), there are some general constraints that extended matter orbiting configurations must abide on the curved spacetimes---\citet{abrafra}. This is an important feature of \textbf{SMBH} accretion disks validating also the relevance of the \textbf{RAD} model. However, accreting disk models are generally associated to specific characteristics of their \textbf{BH} attractors and especially their mass range\citep{abrafra}. The geometrically thick disks, which were adopted here as aggregate components, are governed by the gravitational forces predominant with respect to other possible components of the force balance, and they are therefore associated to \textbf{SMBHs} spacetimes where the curvature effects and the fluid rotation are significant in the determination of the toroidal equilibrium and morphology. \subsection{Phenomenology} \textbf{RAD} phenomenology is significant for the high energy phenomena related to accretion onto \textbf{SMBHs} in \textbf{AGNs} which could be observed in their X-ray emission. \textbf{RAD} framework also includes the jet launch with an interesting and intricate shell structure. These structures then fit into the more broad discussion on the role and significance of open surfaces in relation to (matter) jets emission and collimation, as well as jet-accretion correlation--see \cite{KJA78,Sadowski:2015jaa,Lasota:2015bii,Lyutikov(2009),Madau(1988),Sikora(1981)}. Among the other phenomenological application of these studies, there is a possible connection between \textbf{RADs} seismology and \textbf{QPOs}--low and high frequency peaks in the power density spectra--see also \citet{Ingram:2016tbq}. The pattern of the possible oscillation modes of the tori aggregate has been provided and related to the evolution of instabilities in \textbf{RAD}%, resulting eventually in combining each aggregate torus %own axis-symmetric and incompressible with the global oscillations of the \textbf{RAD} considered in \citet{ringed,2013A&A...552A..10S}. Missions like \textbf{XMM-Newton} or \textbf{RXTE}, \textbf{NuSTAR} (\textbf{Nu}clear \textbf{S}pectroscopic \textbf{T}elescope \textbf{Ar}ray)\footnote{https://www.nustar.caltech.edu/} are on the verge of these studies--see also \citet{Mossoux:2014fqa,Gilli:2006zi} and \citet{Gandhi:2017dix,Masini:2016yhl,Harrison:2013md}. The current status of observational mission and the potentialities in the next future provide certainly an increasingly high degree of details with which we can look at \textbf{BHs} and accretion disk morphology. The results outlined here thus are directly comparable with the current data and encouraging us to apply already these results for a re-interpretation of the current data analysis set up which is almost unanimously based on the scenario of \textbf{BH} one-accretion disk system -- see for example \citet{Hamb,Hamb1}. Possible evidence of the existence of the ringed accretion disk can be inferred from the study of the optical properties of the ringed-like structures \citep{KS10,S11etal,made,made0}. {Concerning the \textbf{RAD} optical appearance we expect that tori emanating radiation would be distorted into a belt-like configuration because of the curvature effects of the Kerr geometry. (Then the Lense--Thirring effect becomes clearly more important for corotating tori approaching the \textbf{BH}). This aspect will be considered in planned future investigation. We expect that optical phenomena could reflect the \textbf{BH} classification as they differentiate between corotating and counter-rotating tori in dependence of the \textbf{BH} dimensionless spin. Optical properties depend also on the geometrical thickness of the disk and the presence of a disk atmosphere (some accretion disks atmospheres can show similarity with the upper main sequence stars, where magnetic fields and winds appear \citep{1967ApJ...148..217W}. The analysis of the optical appearance of the thick accretion disk is also based on the model of hot coronae above the surfaces of accretion disks and a hot coronal layer implied in the emission profile-- \citep{dance,FW07,made0,made,Schee:2008fc}.} \textbf{BH} \textbf{RAD} model may be revealed by future X-ray spectroscopy, from the study of excesses on the shape of the relativistically broadened spectral line profile, related to a sort of rings model which may be adapted as a special case of the \textbf{RADs}. Specifically, in \citet{KS10} extremal energy shifts of radiation from a ring near a rotating \textbf{BH} were particularly studied: radiation from a narrow circular ring shows a double-horn profile with photons having energy around the maximum or minimum of the range (see also \citet{Schee:2008fc})\footnote{ Notice that significant influence of the self-occultation effect on the profiled spectral lines was for the first time demonstrated in \cite{1992ApJ...400..163B}.}. This energy span of spectral lines is a function of the observer's viewing angle, the \textbf{BH} spin and the ring radius. % Accordingly, the ringed disks may be revealed thought detailed spectroscopy of the spectral line wings, claiming for observation and data re-analysis in this framework, especially in relation with the \textbf{BH} classes considered here. } \subsection{RAD phenomenology and future perspectives for the RAD extensions} On one side, the possibility of more orbiting tori opens new perspectives of enriched accretion phenomenology connected to the \textbf{RAD} structure and its own internal dynamics. \textbf{RAD} typical effects as the double accretion phase or the presence of screening tori which are located between the \textbf{BH} and the accreting torus, or tori between two accreting \textbf{RAD} tori, the presence of \textbf{AGN} obscuring torus, are proved to be expected in specific contexts and under precise conditions on the tori and the central \textbf{BH} attractor. In this analysis we present limits indicating the parameter ranges where these situations occur. On the other side, in the \textbf{RAD} frame, we pose rather narrow constraints for different ad hoc models of several accretion processes which are currently considered, ruling out many of the assumptions considered so far on screening and obscuring phenomena. Those limits and some main indications on \textbf{RAD} observation are given below. This work provides a new scenario for the data interpretation which has been until now essentially framed into the single axi-symmetric, corotating, accretion disk model. Screening and obscuring tori in accretion processes are essentially positioned ad hoc in order to fit the observations and models. It is very likely that at least in the initial phases of the \textbf{RAD} formation, ringed structures may appear, and we explicitly call for a review of the analyses carried out so far in the simplest scenario of one torus model, shifting this classic setup to the model of orbiting \textbf{RAD} axi-symmetric structures including the counterrotating tori. Indeed, the predictions implied by the adoption of a \textbf{RAD} system fit some features of the accretion disk physics which puzzle the theoretical grounds of one class of accreting model, for example the geometrically thick tori with some features proper of geometrically thin disks. A further important aspect in this regard stands in the fact that a \textbf{RAD} could have been disguised so far as one geometrically thin disk; in \cite{ringed} it has been proved that \textbf{RAD} can be treated as a geometrically thin disk blending with distinctive features of geometrically thick torus as super-Eddington luminosity (high accretion rates). This situation can fit a particular hypothesis made to justify the masses of \textbf{SMBHs} located in high redshift galaxies. Concerning the possibility of screening tori we proved that such tori \emph{must} be corotating and screening effects could be possible \emph{only} if the outer accreting disk is counterrotating, but this in fact rules out several of the models currently considered in the screened accretion. According to our analysis screening effects are likely to be observed around high spin \textbf{BHs}--precise limits and constraints are in Table\il(\ref{Table:nature-Att}). Further constraints on thickness and extension on the equatorial plane of the \textbf{RAD}, relevant to set power spectra obscuration follow, from more detailed analysis in Table (\ref{Table:nature-Att}). Those constraints will emerge as distinctive features of the emission spectra. Most interestingly a screening effect may happen in the occurrence of a double accretion phase as in Fig.\il(\ref{Fig:but-s-sou})-(a). The possibility of a double accretion from an $\ell$counterrotating couple is a new feature firstly presented in the \textbf{RAD}, {we clamor for attention on the scientific community to focus on observations on high masses and high spin Kerr black hole where pieces of evidence of this phenomenon should be found}. Besides we proved that no more then two accreting tori can be observed. The detection of a double accretion phase (associated to a double shell of jet emission) would imply the presence of an outer counterrotating accreting torus and an inner corotating one, the central \textbf{BH} is not screened. Vice versa, a quiescent corotating torus, screening the outer retrograde accretion, and the inner prograde one are possible. Such inert screening tori must be relatively small configurations. The analysis of the \textbf{RAD} characteristics which could be extracted from the spectrum emission or accretion rates/luminosity analysis, \cite{long}, may be used to locate the \textbf{RAD} in one of the attractor class of Table\il(\ref{Table:nature-Att}), thus identifying the central \textbf{BH} attractor. Spectra, for example from X-ray emission, would show evidence of the \textbf{RAD} spacing. Note that the distance in spin boundaries of the \textbf{SMBHs} classes can be very close, also up to $10^{-2}M$, this means that we can distinguish, considering \textbf{RAD} features, a spin range $\Delta a\approx 10^{-2}M$, depending on the constraints provided on the accretions properties. It should be noted that a large part of the accretion disk analysis is actually dealing mainly with corotating disks, our \textbf{RAD} analysis provides more precise insight on the physics connected with counterrotating tori. Screening effects are possible only in the case of accreting counterrotating tori, this implies that the current analysis using a screening torus must take into account the fact that the outer torus, in the \textbf{BH}-screening torus-accreting torus system, is counterrotating. This is a strong restriction in the analysis of retrograde disks. On the other hand, if there is a system with a screening torus and an obscuring one, this situation implies that the \textbf{BH} spin has to be relatively small $a\lessapprox 0.46M$. Most importantly a corotating torus \emph{cannot} be screened to the \textbf{BH}, i.e., no screening effect from axi-symmetric inner torus can be set--this means that many screening models are effectively ruled out. A further interesting aspect is constituted by jet emission in \textbf{RADs}. \textbf{RADs} are related to jet emission in two different ways. Firstly, each toroid considered as \textbf{RAD} component allows open funnels, proto-jets, occurring when the centrifugal component in the disk force balance is strong enough to push matter along the vertical direction along the \textbf{BH} axis. Jet for retrograde tori is another interesting feature. We considered these open solutions too in the set-up of the attractor classifications, although a focused investigation of these are especially in \cite{open,long,app}. Secondly, as jets are expected to be correlated to accretion and especially to the inner part of an accreting disk, the dynamical picture featuring more accreting tori, up to maximum of two $\ell$counterrotating accreting tori, including the possibility of screening and obscuring tori, is certainly an interesting perspective in the context of jet emission. We proved that there can be up to two jets (no proto-jets), i.e., double shell of jets. According to constraints on the double accreting phase, the outer jet being correlated to the outer retrograde \textbf{RAD} torus, and the inner jet from the inner corotating, accreting torus. Constraints on spacing and tori dimensions are set according to \textbf{SMBH} spin and details on these constraints are presented in Table\il(\ref{Table:nature-Att}). Spacing between the tori will reflect in the jets separations. The change of \textbf{BH} spins will reflect in change of the spacetime and other features of jet emission. In this context, it is proper to stress that characteristic \textbf{RAD} features in accreting processes, such as screening and obscuring tori act also in regard of the jet emission, the presence of inert materials will affect the energy release, the jet launching and collimation. Small corotating tori will be located eventually between two accreting tori and inside the double jet shall, while in the case of one jet correlated to a counterrotating accreting torus, the \textbf{BH}-jet system, can be possibly characterized by an inner, screening, corotating torus. It is then crucial to establish possible \textbf{RAD} evolutions following the occurrence of the instability processes in the structure. Constraints on \textbf{RAD} evolution have been discussed more extensively in \cite{dsystem,long}. Generally, the situation after tori collision emergence, or accretion between two tori, with matter impacting from one torus to another, can result also in the destruction of the \textbf{RAD}, i.e. formation of one torus after tori merging. Evolution paths in these models were discussed in \cite{dsystem}, energy release associated to the collision and accretion phases including the evaluation of the torus cusp luminosity and accretion rates are examined in \cite{long}, where evolution is also discussed in context of energy release in \textbf{RAD}. We can say that one possibility is the merging after accretion of the outer counterrotating torus of a couple or collision with one inert torus or during the double accretion phase; in all these cases there will be a modification of the stability properties of each involved torus, according to a variation of the parameters range, for matter supply from the outer accreting counterrotating torus, followed by an alteration of the fluid specific angular momentum and an increase of the $K$-parameter, reflecting change in density and tori dimension. Such phenomenon has to be considered, focusing on a more accurate prescription of the torus inner dynamics. The possibility that the accretion or collision may in fact not lead to the destruction of the \textbf{RAD} inner structure, but rather to a sequence of instability processes (accretion phases) involving an inner torus has also been discussed. Concerning \textbf{RAD} evolution, it must be said that the oscillation modes of each toroidal component, will be combined together in the ringed structures to reflect the QPOs phenomenon--see discussion in \cite{ringed} for the \textbf{RADs} perturbation. One can assume oscillations of an individual torus of the \textbf{RAD} system, and possible excitation of tori oscillation by external influence, e.g, due to original oscillation of the first oscillating torus. The resonance models of the epicyclic frequencies could give relevant explanation of double peak \textbf{QPOs}. Such double \textbf{QPOs} are recently directly observed in \textbf{AGNs} (see \cite{Carpano}) and they might be directly linked to the \textbf{RAD} systems, giving independent estimates of the mass and spin of the Kerr \textbf{SMBH}. {Finally, there are possible relevant extensions of the \textbf{RAD} model to be considered, according to the specific frame where the aggregate hypothesis may be applied. There are two main classes of model generalizations, the first involving a change in each aggregate component, for example in \cite{Fi-Ringed} magnetized \textbf{RADs} have been considered. Secondly, assumptions of \textbf{RAD} entire structure may be modified, particularly on regards of the aggregate symmetries. We discussed the relevance of tori misalignment in \textbf{RADs} in several parts of this work. Considering the different properties of in-falling matter from several companions during the accretion processes into the Kerr black hole, a tori misalignment is highly expected at least in the early phases of formation of the \textbf{RADs}. The presence of initial misalignment will lead also to a change of the entire \textbf{BH}-disk system inducing a variation of the \textbf{BHs} properties due to modifications of the spin magnitude and orientation. The \textbf{BH}-\textbf{RAD} tori considered here might be seen as the final state of this complex scenario. In fact in several cases the final ending of the \textbf{BH}- misaligned tori dynamics results in equatorial disks. %Assuming then the first \textbf{RAD} model to be constituted by equatorial and axes-symmetric tori has lead to a first set-up for further analysis where different situations will be considered as deviation with respect to the symmetries as well as and the matter structure. {Wind and jet emission can be present in \textbf{RADs}, showing a very complicated behavior. This might possibly lead to non linear dynamics, and as a consequence of this, the \textbf{RAD} system may be subjected to deterministic chaos (e.g. \cite{May}). Moreover, toroidal components of the \textbf{RAD} can exchange matter and angular momentum leading to different effects, this situation has been discussed in an evolutionary scenario in \cite{dsystem}: a possibility includes ``drying-feeding'' effects or { evolutionary loops among tori} of a \textbf{RAD} couple as described in \cite{dsystem}. Another possibility in this two-steps exchange of angular momentum between tori, is the excitation of oscillations which could eventually be observed as it occurs in the matter and angular momentum exchanges in certain class of binary stars. This behavior in general could be connected with the quasi-periodic or chaotic oscillations of the \textbf{RADs}. We shall deepen this aspect of the \textbf{RAD} dynamics and the model generalizations in future works.} % Particularly inclined tori will be set first in the simplest case of the spherically symmetric background of the Schwarzschild black hole and then slow rotating spacetimes will be analyzed. The ranges of spin variation in the slow regime will be located considering also the limits provided in the \textbf{SMBHs} classes provided of Table\il(\ref{Table:nature-Att}). A further extension on the model affecting tori symmetries is for the warped disk analysis where the tori warping will be included.} {Then in binary \textbf{BHs} systems (\textbf{BBHs}), \textbf{BHs} spins are usually not aligned \citep{Moran:2008dv}, a complex \textbf{RAD} related problem is considering \textbf{BBHs} with two mini-disks and a circumbinary disks, with not aligned \textbf{BHs} spins.} % | 18 | 8 | 1808.00238 |
1808 | 1808.03834_arXiv.txt | { Although the globular clusters in the Milky Way have been studied for a long time, a significant fraction of them lack homogeneous metallicity and radial velocity measurements. In an earlier paper we presented the first part of a project to obtain metallicities and radial velocities of Galactic globular clusters from multiobject spectroscopy of their member stars using the ESO Very Large Telescope. In this paper we add metallicities and radial velocities for a new sample of 28 globular clusters, including in particular globular clusters in the MW halo and the Galactic bulge. Together with our previous results, this study brings the number of globular clusters with homogeneous measurements to $\sim {69}$\% of those listed in the W. Harris' catalogue. As in our previous work, we have used the \ion{Ca}{ii} triplet lines to derive metallicities and radial velocities. For most of the clusters in this study, this is the first analysis based on spectroscopy of individual member stars. The metallicities derived from the \ion{Ca}{ii} triplet are then compared to the results of our parallel study based on spectral fitting in the optical region and the implications for different calibrations of the \ion{Ca}{ii} triplet line strengths are discussed. {We also comment} on some interesting clusters and investigate the presence of an abundance spread in the globular clusters here. A hint of a possible intrinsic spread is found for NGC\,6256, which therefore appears to be a good candidate {for further study}. } | Globular clusters (GCs) were considered as simple stellar populations for a long time, with most of their parameters derived from heterogeneous methods and data, largely driven by the technology available at the time when the clusters were studied. {Recently,} with the acquisition of extremely precise photometry from Hubble Space Telescope (HST) imagers and spectroscopy from large telescopes, the idea of simplicity has been replaced by a more complex view of their star forming history \citep[see, e.g.,][]{gratton+2012}. {Closely linked to that history is the dynamical evolution of clusters, which is another area of study that benefits from homogeneous samples and reduction techniques. Data sets like these can be used as input to N-body simulations, and the resulting dynamical parameters allow comparative studies such as that recently published in \citet{baumgardt+2018}.} Improvements in our understanding of GC evolution came also from the introduction of large spectroscopic studies, which enabled measuring homogeneous abundances for large samples of GC stars. The largest sample of metallicities measured with the same method comes from the study of equivalent widths of the \ion{Ca}{II} triplet (CaT) lines in the spectra of GC stars by \citet{rutledge+1997} who determined metallicities for 52 clusters and made a literature compilation on their scale for a total of 71 clusters. The second largest sample is represented by integrated-light spectroscopic studies such as those of \citet[][ZW84]{zinn+1984} and \citet{armandroff+1988} which account for an additional $\sim20\%$ of the metallicity measurements for GCs. Another approach to metallicity homogenisation was followed by \citet{carretta+2009} using their high-resolution spectroscopic analysis of 19 GCs to set a new metallicity scale based on \ion{Fe}{i} lines. Transformations to the new scale were given for 114 GCs from \citet{harris+1996}. The lack of homogeneous measurements currently mostly affects the outer-halo (distant) and bulge (highly extincted) globular clusters, which give important constraints on the Milky Way formation. This motivated our group {to increase the sample of GCs with metallicities determined in a consistent way. The project is based on medium-resolution spectroscopy of individual red giant branch (RGB) stars in a sample of clusters which complements that of \citet{rutledge+1997}. The spectra were obtained with FORS2 at the ESO VLT observatory in two spectral intervals, the CaT region and the green region that is rich in metal absorption lines.} In Paper~I \citep[][S12]{saviane+2012} we presented the methods and analysis of spectra in the CaT region for a first set of clusters from the original sample, including eight {calibration} GCs and 20 programme clusters, four of which are in common with \citet{rutledge+1997}. {The use of spectra at the CaT} is particularly useful for targets near the Galactic plane or towards the Milky Way bulge because of the lower extinction in the far-red spectral range. A new CaT metallicity calibration based on the \citet[][hereafter C09]{carretta+2009} abundance scale was derived and used to measure the metallicity of the programme clusters. In addition, the clusters from \citet{rutledge+1997} were included, converted to the new scale. In a parallel project, new metallicity measurements were obtained for a largely overlapping GC sample from {full spectral fitting} of medium resolution spectra in the visible region \citep{dias+2015,dias+2016}. The new metallicity scale, independent of previous empirical calibrations, appears to be consistent both with abundances from high-resolution spectroscopy and photometric metallicities \citep{cohen+17}. On the basis of the new scale, and taking advantage of an increased sample of metal-rich clusters, \citet{dias+2016eso} {generated a homogeneous compilation} of metallicities for 152 GCs (97\% of the total) from different sources. This paper increases our homogeneous database by adding new CaT measurements for a further set of 28 GCs, which in combination with those published in \citet{saviane+2012} provide metallicity estimates for ${\sim 69\%}$ of the Milky Way globular clusters in the \citet{harris+1996}\footnote{Hereafter all references to the catalogue refer to the 2010 on-line version \citep[][]{harris+2010}, unless otherwise specified.} catalogue. The [Fe/H] estimates from the CaT are then compared to the GC metallicities independently obtained from the green spectral interval \citep{dias+2015,dias+2016}. This paper is organised as follows: in Sect.~2 the observations and data reduction are presented; in Sect.~3, radial velocities are measured and compared with literature values; in Sect.~4 several CaT EW calibration relations are discussed and metallicities are estimated. Section~5 provides comments for some interesting clusters. Finally, our results are summarised in Sect.~6. \begin{figure}[t] \centering \includegraphics[angle=0,width=1.0\columnwidth]{svasquez_1.eps} \caption{Reddening and visual apparent magnitude at the level of the horizontal branch for the globular clusters observed in \citet{saviane+2012} and this work. Data are from the Harris catalogue. The dashed line shows the $V$ magnitude of the horizontal branch (assuming $V_{0,\rm{HB}}=0.5$) for a cluster at the distance of the bulge (8~kpc) and subject to different degrees of extinction.} \label{reddening_plot} \end{figure} | \begin{figure}[t] \centering {\includegraphics[width=1.0\columnwidth]{svasquez_9.eps}} \caption{Metallicity distribution of the MW GCs in the merged CaT sample. The thick line shows the histogram made with our data only (Paper~I and the present work); the remaining data are from \citet{rutledge+1997} as put on the C09 scale by \citet{saviane+2012}. The error bars are Poisson errors on the number of clusters in each bin.} \label{fhisto} \end{figure} We have obtained new radial velocities and metallicities for stars in 28 Milky Way GCs, {as the completion of our spectroscopic survey} aimed at providing a homogeneous set of metallicities from medium resolution data in two spectral regions. The results in this paper are based on FORS2 spectroscopy with a resolution $R\sim2500$ in the region centred on the \ion{Ca}{II} triplet lines ($8450-8700$ \AA) {of individual cluster candidate red giants.} In a previous study of this series, \citet{saviane+2012} measured the metallicity of 20 programme clusters using eight template clusters with known metallicity to establish an empirical CaT line strength calibration. In this paper we have measured the metallicity of an additional sample of 28 globular clusters mainly in the Galactic bulge or far in the halo. Most of the clusters analysed here are therefore faint and/or highly reddened, which for a long time prevented any metallicity measurement based on spectroscopy of individual stars. {We provide metallicities based} on the CaT calibration of \citet{saviane+2012}, with typical uncertainties of the order $\sim0.15$ dex. We also discussed the impact on the measured [Fe/H] of alternative empirical calibrations such as those of \citet{vasquez+2015} and \citet{dacosta+2016}, with differences especially noticeable in the metal-rich regime (${\rm [Fe/H]>-0.8}$). {Our sample is complemented by the 71 GCs published by \citet{rutledge+1997} transformed to our metallicity scale} \citep{saviane+2012}. Altogether, our project provides the largest GC sample with homogeneous metallicity measurements from CaT currently available. The combination of all data yields CaT metallicity estimates on the same scale for {107 GCs, that is ${{\sim69\%}}$} of the clusters published in H10. The metallicity distribution of all clusters in our CaT project is plotted in Fig.~\ref{fhisto}, where the new data represent about 50\% of the available measurements for metal-rich clusters. As a test of the metallicity scale provided by the \ion{Ca}{ii} triplet, we compared our results with those obtained from spectroscopy of stars in the same clusters using a spectral fitting technique in the optical region \citep{dias+2015,dias+2016}. The comparison shows that the metallicities measured from the CaT line strengths using the calibration of \citet{saviane+2012} are consistent within the errors with the [Fe/H] values derived from spectral fitting by Dias et al., which in turn agree with metal abundances {from high-resolution spectroscopy.} Our study confirms the CaT method as a powerful tool to determine [Fe/H] for stellar systems so distant or reddened that high-resolution spectroscopic analysis becomes very difficult or prohibitive, provided that the calibrators are selected from the same population as the targets. {This will be particularly useful to study resolved stars beyond the Local Group with the new generation of extremely large telescopes.} We have identified some interesting cases among the 28 GCs analysed in the present work. We obtained a much improved metallicity measurement for Djorg\, 2: while previous work based on a NIR CMD classified it as a metal-rich globular cluster with a metallicity ${\rm [Fe/H]=-0.65}$, our EW measurements and optical photometry suggest a more metal-poor value ${\rm [Fe/H]\approx-1.1}$. Additionally, we found a radial velocity $v_{\rm r}=-160$ km s$^{-1}$ which is the first estimate for this cluster. Another interesting case is that of Terzan 1, for which we found a difference $\Delta v_{\rm r}=-51$ km s$^{-1}$ in radial velocity with respect to previous values. Our estimate is consistent with the latest measurement from high-resolution spectra \citep{valenti+2015}. Finally, as in \citet{saviane+2012}, we explored the presence of intrinsic metallicity dispersions within our sample: we found three clusters that might have a metallicity dispersion, based on a comparison of their EW dispersion with the intrinsic error in EW measurements (parameter $\rho\geqslant1.3$). One of them, NGC\,6256, is the most probable candidate, with $\rho=2.2$. For this cluster our data suggest an intrinsic metallicity dispersion $\sigma_{\rm [Fe/H]}=0.2$ dex, but more data are needed to confirm this hint. | 18 | 8 | 1808.03834 |
1808 | 1808.08251_arXiv.txt | {Modeling emission lines from the millimeter to the UV and producing synthetic spectra is crucial for a good understanding of observations, yet it is an art filled with hazards. This is the proceedings of \quotes{Walking the Line}, a 3-day conference held in 2018 that brought together scientists working on different aspects of emission line simulations, in order to share knowledge and discuss the methodology. Emission lines across the spectrum from the millimeter to the UV were discussed, with most of the focus on the interstellar medium, but also some topics on the circumgalactic medium. The most important quality of a useful model is a good synergy with observations and experiments. Challenges in simulating line emission are identified, some of which are already being worked upon, and others that must be addressed in the future for models to agree with observations. Recent advances in several areas aiming at achieving that synergy are summarized here, from micro-physical to galactic and circum-galactic scale. } \keyword{simulation; line emission; galaxies; ISM; radiative transfer; hydrodynamic simulations; CGM; AGN} \begin{document} | Line emission from the Interstellar Medium (ISM) of galaxies carries information that is crucial in understanding galaxy evolution. Observing line emission across the electromagnetic spectrum allows us to characterize the mass, composition, and chemical state of the ISM, as well as to trace galaxy properties such as star formation rate (SFR), metallicity and dynamics. For example, the emission from major cooling lines, such as H$\alpha$ or \cii, is sensitive to the physical conditions (densities, radiation field) and dynamics of the ISM. In addition, emission lines work on all physical scales, from galaxy dynamics and inflows to turbulent and collapse motions in star-forming clouds and cores. By systematically comparing spectral-line signatures of different physical models, one can correctly identify the physical processes occurring in these regions. Furthermore, the emission from ionized interstellar gas contains particularly valuable information about the nature of the ionizing radiation sources in a galaxy. In fact, prominent optical emission lines are routinely used to estimate whether ionization is dominated by young massive stars (tracing SFR), an AGN or evolved, post-asymptotic giant branch (post-AGB) stars. Three of the most widely used line-ratio diagnostic ``BPT'' diagrams\footnote{``Baldwin, Phillips \& Terlevich'' (BPT) diagrams \cite{Baldwin1981}}, relate the [OIII]/H$\beta$ ratio to the [NII]/H$\alpha$, [SII]/H$\alpha$ and [OI]/H$\alpha$ ratios. These diagrams have proven useful in identifying the nature of the ionizing radiation in large samples of galaxies in the local Universe \citep{Kewley2001,Kauffmann2003}. Complementary to line emission are the observations of absorption lines of the circumgalactic medium (CGM), which can give key information on the history of the feedback, in terms of chemical, ionization, and thermodynamical state of the outflowing/inflowing gas, that regulates the star formation process. Gas kinematics, from both emission and absorption, give information about large scale gas flows. Thus galactic outflows, from active galactic nuclei (AGN) and starbursts, can be combined with CGM absorption line observations, to study the star formation history, AGN activity history, and feedback processes that regulate both the evolution of the galaxy and its environment. Looking back on the past three decades, the approach to creating synthetic observations of line emission has gone from simplified analytical modeling to complex simulations, increasing the reliability of the results. Since many important cooling lines emerge from the photodissociation regions (PDRs) of the ISM, modeling line emission from the PDRs has been an active field of research since the basic 1D PDR models came into place in the 80’s \citep[e.g.][]{Tielens1985,vanDishoeck1986,vanDishoeck1988a,vanDishoeck1988b,Sternberg1989}. This early modeling work included the reprocessing of starlight in the UV to infrared continuum by dust grains and polycyclic aromatic hydrocarbons (PAHs), and the stratified layers of increasingly photo-dissociated species as one moves through the neutral gas of a cloud towards the HII region. This picture is illustrated in the top panel of Fig.\,\ref{fig:tielens85}. In the late 90s, the models presented in the 80's \citep{Tielens1985,Hollenbach1991} were used as basis for a modeling effort to create line ratio diagnostic plots of the \oi, \cii, \ci and CO FIR emission lines \cite{Kaufman1999}. Later the modeling was improved with the online PDR Toolbox\footnote{\url{http://dustem.astro.umd.edu/pdrt/}} as a result \citep{Pound2008,kaufman2006}. \begin{figure}[H] \centering \includegraphics[width=0.6\linewidth]{tielens85.png}\\ \includegraphics[width=0.495\linewidth]{pdr_densities.pdf} \includegraphics[width=0.495\linewidth]{pdr_cooling.pdf} \caption{{\it (Top)} Sketch of standard PDR structure, from innovative work in the 80's \citep[][\copyright AAS. Reproduced with permission]{Tielens1985} . {\it (Bottom)} For comparison, an example of PDR structure computed this year with version C17 of \cloudy \cite{Ferland2017} for model V1 of \cite{Rollig2007}, adapted to extend to $A_V = 100$. {\it (Left panel)} Densities of important atoms and molecules in the PDR. {\it (Right panel)} Local emissivities for some of the most important coolants.} \label{fig:tielens85} \end{figure} Still, one of the main problems when comparing these models to actual observations is that the models are based on patches of gas with constant density, metallicity, and radiation field strength whereas any given galaxy (or region in a galaxy) will be a superposition of different gas states and radiation conditions as illustrated in Fig.\,\ref{fig:superimposed_clouds} (Ideally, each cloud should also have a radial gradient, but this feature has been omitted from the Figure for the sake of clarity). With improved simulations of galaxy (and cloud) formation, it became possible in the 21st century to use galaxy/cloud simulations as direct input to photoionization codes, that also calculate the line emission \citep[e.g.][]{Bolatto1999,Rollig2006,Narayanan2006,Popping2014,Vallini2015,Olsen2015,Popping2016,popping2018,vallini18}. This new approach gave way for a more realistic picture in which each region of the ISM is approximated by a combination of different gas conditions. Some of these simulations include radiative transfer and the non-equilibrium impact on the gas temperature of the local radiation field generated by nearby star formation \citep[e.g.][]{Rosdahl2018}, allowing for more accurate estimation of the nebular line emission. This method is being extended to simulations of the CGM, with increased resolution to capture the impact of outflows on accreting structure (see Sec.\,\ref{sec:44}). \begin{figure}[H] \centering \includegraphics[width=0.55\linewidth]{superimposed_clouds.png} \caption{Illustration of the difference between simulating line emission from a cell of constant density, temperature, metallicity etc., and using an ensemble of clouds superimposed to create a more realistic synthetic observation of a patch of the sky. } \label{fig:superimposed_clouds} \end{figure} Improved numerical techniques such as those mentioned above have also meant a slight division in the field of line modeling, between groups would who specialize in creating the photoionization codes and groups that apply these codes to hydrodynamical simulations. The photoiononization codes have evolved to also calculate line emission and are kept updated as experimental values for collision strengths, chemical rate coefficients and photoelectric heating rates are being revised and databases of atomic and molecular line data increase in size \citep[e.g.][]{Keto+04,Brinch2010,Keto_Rybicky10,Dullemond2012,Yajima2012,Ferland2013,Krumholz2014,Gray2015}. Similar to the case of PDRs, significant progress has been achieved over the last few years in modelling nebular emission from ionized regions around young star clusters, AGN and post-AGN stellar populations in a full cosmological context. Yet, fully self-consistent models of this kind are currently limited by the performance of cosmological radiation-hydrodynamic simulations and insufficient spatial resolution on the scales of individual ionized regions around stars and active nuclei. To circumvent these limitations, some pioneer studies proposed the post-processing of cosmological hydrodynamic simulations and semi-analytic models with photoionization models to compute the cosmic evolution of nebular emission \citep{Kewley2013,Orsi2014,Shimizu2016}. Some further improvement has been achieved recently by \cite{hirschman2017}, not only accounting for the integrated nebular emission from young stars \citep[as in][]{Orsi2014,Shimizu2016}, but also from AGN and post-AGB stars, based on the star formation, AGN luminosity, gas density, and chemical enrichment histories of the simulated galaxies. Across different aspects of simulating line emission, key questions remain to be answered: \begin{itemize} \item What are the best emission lines to trace various ISM properties and ionizing sources in galaxies? \item How can we use emission lines to trace feedback and ISM evolution with redshift? \item How should absorption features be correctly interpreted? \item Where do we stand in deriving sub-grid physics and comparing codes? \item How do we coordinate our efforts? \end{itemize} To address such questions, the conference \quotes{Walking the Line} was held in Phoenix, Arizona on March 14-16 2018. The main topic was Simulating Line Emission from Galaxies\footnote{\url{https://walk2018.weebly.com/}} and the program\footnote{\url{https://walk2018.weebly.com/program.html}} included 35\,min presentations from three invited speakers and 15\,min presentations from 24 of the remaining 27 participants. In total, nine participants were PhD students and 1/3 of all participants were women. Most talks are now publicly available with video recordings online\footnote{Slides from and video recordings of talks can be found at \url{https://zenodo.org/communities/walk2018/}}. In addition, the conference had daily discussion sessions stimulating in-depth exploration of new topics and sparking new connections. This way, the meeting was an opportunity to discuss the challenges that need to be solved in order to make more realistic simulations and a fairer comparison with observations. To take a specific case study where line emission has come to play an important role, one can consider the problem of measuring gas mass, which became a common thread for a large part of the workshop. Measuring the gas mass of galaxies, the fuel for star formation, as a function of cosmic time is a crucial component in understanding the formation and evolution of galaxies. Different groups have observed the gas content of main-sequence galaxies through cosmic time (the bulk of the galaxy population that is responsible for forming new stars at any epoch in cosmic history) either through their $^{12}$CO (hereafter CO) emission or the dust continuum \citep[e.g.,][]{Aravena2010,Daddi2010,Tacconi2010, Tacconi2013, Tacconi2018,Geach2011,Magdis2012,Santini2014,Bethermin2016,Decarli2016,Scoville2016}. Both approaches (CO and dust) rely on uncertain conversion factors between the observed luminosity and an estimated gas mass. ISM physical processes can further complicate the use of CO and dust continuum to reliably measure the gas mass of galaxies. For example, the contrast of the CO emission line and dust continuum against the cosmic microwave background becomes lower at higher redshifts \citep[e.g.,][]{daCunha2013,vallini18} and under the influence of cosmic rays the CO molecule can be destroyed \citep{Bisbas2015,Glover2016}. Reliable alternative measures of the gas mass in a well defined sample of galaxies are therefore essential to overcome the systematic uncertainties in the CO and dust conversion factors and to overcome the CMB contrast. Alternative options seen in the literature include for instance the emission from [CI] \citep{Bothwell2016,Popping2017}, [CII] (Zanella et al. in prep, 2018), from H$_{\rm 2}$O (e.g. Liu et al. 2017) and PAHs features (e.g. Cortzen et al. submitted 2018), and from optically thin isotopologues \citep[e.g.,][]{cormier2018}. The synergy between galaxy formation simulations, ISM chemistry simulations, and radiative transfer codes has the power to pave the way towards reliably measuring gas masses. This synergy allows for a controlled setting in which the conversion between sub-mm line/continuum strength and H$_{\rm 2}$ mass can be explored. This is important to 1) quantify under which physical conditions classical approaches to measure gas masses such as CO and dust continuum emission break down and 2) explore the robustness of other mid-IR and sub-mm emission lines such as PAHs, [CI], [CII] and H$_{\rm 2}$O (although one could think of other examples) as a tracer of molecular hydrogen mass. In this paper we summarize the topics covered in the conference, with particular attention to the break-out sessions of this workshop. We report our conclusions on the state of the art of modelling emission lines, by going to progressively larger physical scales (Sec.s \ref{sec:2} -- \ref{sec:4}), concluding by discussing the possible ways to move forward as a community (Sec. \ref{discussion}). \begin{table}[H] \caption{Overview of the lines discussed at the workshop and referred to in the remainder of this paper.} \label{table:2} \centering \begin{threeparttable} \begin{tabular}{p{1cm}|p{2cm}|p{3cm}|p{4cm}|p{2cm}} \toprule \textbf{Name} & Type & \textbf{Wavelength(s)\tnote{(1)}} & \textbf{Tracer of} & \textbf{Reference for wavelength(s)}\\ Ly$\alpha$ & Recombination & $1215.67\,$\AA & Ionized ISM & \cite{Sansonetti2004} \\ C\,{\sc iv} & CE\tnote{(2)} & $1548.19$, $1550.77\,$\AA & Stellar wind, ionized ISM & \cite{Leitherer2011} \\ O\,{\sc iii}] & CE\tnote{(2)} & $1660.81$, $1666.15\,$\AA & Ionized ISM & \cite{Leitherer2011} \\ He\,{\sc ii} & Recombination & $1640.42\,$\AA & Stellar wind, ionized ISM & \cite{Leitherer2011} \\ $[$C\,{\sc iii}$]$ & CE\tnote{(2)} & $1906.68$ \AA & ISM & \cite{Leitherer2011} \\ C\,{\,\sc iii}] & CE\tnote{(2)} & $1908.73\,$\AA & Ionized ISM & \cite{Leitherer2011} \\ H$\beta$ & Recombination & $4861.36\,$\AA & Ionized ISM & NIST\tnote{(3)} \\ $[$O\,{\sc iii}$]$ & CE\tnote{(2)} & $4958.91,5006.84$ \AA & Ionized ISM & NIST\tnote{(3)} \\ $[$O\,{\sc i}$]$ & Recombination & $6300.30\,$\AA & Ionized ISM & NIST\tnote{(3)} \\ H$\alpha$ & Recombination & $6562.80\,$\AA & Ionized ISM & NIST\tnote{(3)} \\ $[$N\,{\sc ii}$]$ & CE\tnote{(2)} & $6548.05,\,6583.45$ \AA & Ionized ISM & NIST\tnote{(3)} \\ $[$S\,{\sc ii}$]$ & CE\tnote{(2)} & $6716.44,\,6730.82$ \AA & Ionized ISM & NIST\tnote{(3)} \\ \ci & Fine-structure & $609.14$, $370.42\,\mu$m & Atomic and molecular gas & LAMDA\tnote{(4)}\\ \cii & Fine-structure & $157.74\,\mu$m & All ISM & LAMDA\tnote{(4)}\\ \oi & Fine-structure & $63.18$, $145.53\,\mu$m & Atomic and molecular gas & LAMDA\tnote{(4)}\\ CO & Rotational & $2.6$, $1.3$, $0.87$\,mm ... & Molecular gas & LAMDA\tnote{(4)}\\ \midrule \bottomrule \end{tabular} \begin{tablenotes} \item[1] We give air wavelengths above 2000 \AA~ and below 20,000 \AA, and vacuum wavelengths otherwise.\vspace{3.5pt} \item[2] Collisionally excited line\vspace{3.5pt} \item[3] \url{https://physics.nist.gov/PhysRefData/ASD/lines_form.html}\vspace{3.5pt} \item[4] \url{http://home.strw.leidenuniv.nl/~moldata/}\vspace{3.5pt} \end{tablenotes} \end{threeparttable} \end{table} Note that the number of emission (and absorption) lines observed from the ISM of galaxies is ever increasing, and we did not attempt to cover them all with this workshop. However, to give the reader a quick overview of the lines that were discussed at the workshop and are commonly used to diagnose the ISM and/or CGM, we present in Table \ref{table:2} a list of lines that will be mentioned in this paper. | \label{con} This paper reviews recent progress and identifies current problems in the field of simulating line emission from the ISM and CGM. Conclusions presented here are a result of discussion sessions at the workshop “Walking the Line” held in Phoenix Arizona on March 14-16 2018. This international meeting was the first of its kind to gather astronomers specializing in the modeling of line emission and brought together 30 scientists from grad student to professor level. Most talks with follow-up questioning sessions have been made available online, but as a more digestible release of knowledge, we have condensed the outcome on the most relevant topics in this paper. Returning to the questions posed in the introduction, below are the main conclusions, focusing on the methods used to give the most reliable answers to each question. {\it What are the best emission lines to trace various ISM properties and ionized sources in galaxies?} The best method to answer this question of course depends on the type and scale of the emitting source in the ISM considered. In order to provide the reader with an overview, though not complete, of tools used to calculate line emission either on-the-fly in a simulation or in post process, we listed the software tools that were discussed at the workshop. One outstanding issue that can affect these tools, mostly when considering line emission in the X-ray but also at longer wavelengths \cite{delzanna18}, is to correct theoretical heavy-element energy levels that do not agree with experimental values. Key to deriving the line emission from any cloud is the UV field irradiating that cloud. However, our knowledge on the ionizing continuum of stars is severely hampered by the observational difficulties and thus currently restricted to theoretical predictions. Direct observations of individual, massive, low-metallicity stars will only be possible with telescopes such as LUVOIR and HabEx. Once the UV field is known, radiative transfer of UV into infrared light is a crucial part of any method used, especially when studying neutral or molecular regions. While still too computationally expensive to run alongside galaxy-scale simulations, work is ongoing to include it more regularly in post-process as part of any line emission simulation. In order to compare simulated galaxies with observed ones, we need to put more thought into the choice of physical properties being compared. For example, a SFR that is modeled as instantaneous can hardly be compared to an observed SFR that is essentially derived as an average over 10 to a few-hundred Myr. Many other parameters such as the effects of an AGN or the inclination of a galaxy can distort an attempt to compare model galaxies with observed ones. Finally, it was concluded that the aim should always be to simulate more than one emission line simultaneously, in order to assure consistency across an the entire galaxy. {\it How can we use emission lines to trace feedback and ISM evolution with redshift?} Feedback from AGN in the form of radiation and outflows has a profound effect on the nebular line emission. As an example, simulations of massive galaxies with and without AGN were presented in which the presence of an AGN could clearly be identified from the [OIII]/H$\beta$ and [NII]/H$\alpha$ emission-line ratios. AGN-driven outflows thus strongly regulates the star formation history, at least for massive galaxies, which controls nebular emission from young stars via the ionization parameter. As an observational example, the high-redshift galaxy MRC 0943-242 is believed to contain an AGN, with line emission (and absorption) revealing ionized ISM in close proximity to the nucleus as well as outflows and high levels of turbulence. In close relation to stellar and AGN feedback, turbulence and shocks are another important source of heating and ionization. Recently developed tools for the treatment of turbulence and shocks were presented to aid the line emission simulations in cases where necessitated. On cloud-scale level, simulating the line emission from collapsing dense cores can reveal whether the molecular cloud is undergoing an ``inside-out’’ or ``outside-in’’ collapse. Work is undergoing to break this degeneracy by making synthetic observations of signature spectral lines for comparison with observations. {\it How should absorption features be correctly interpreted?} Simulations of line absorption through the CGM are faced by its own set of problem, where for example the treatment of the ionizing field from the galaxy source can play a crucial role but has yet to be fully treated in a non-equilibrium manner. Another choice to be made when simulating the CGM is the level of spatial resolution needed, as better resolution leads generally to more gas clumps and a point of convergence must be found. {\it Where do we stand in deriving sub-grid physics and comparing codes?} When studying ISM evolution with redshift, line emission must be derived from cosmological simulations which capture the evolution but comes at low spatial resolution compared to semi-analytical models. In addition, full chemistry and radiative transfer must typically be applied in post process. The uncertainties in the approach could be alleviated by more work on scales between galaxies and clouds, because simulations of the latter come at higher resolution and could hence be used as a benchmark on larger scales. {\it How do we coordinate our efforts?} We hope to have set an example for more theoretical workshops of this kind to come, and look forward to seeing the results of the ongoing projects presented here. The developers of the software tools used to derive line emission, increasingly encourage their users to report problems or suggest improvements to their code via online platforms. With the list of caveats identified in this workshop, albeit long, we continue to believe that the future of simulating line emission is bright, especially when research groups in the field come together to find and solve common problems. | 18 | 8 | 1808.08251 |
1808 | 1808.01146_arXiv.txt | We report the discovery of Serenity-18, a galaxy at $z\simeq5.939$ for which we could measure the content of molecular gas, $M({\rm H}_2)\simeq 5 \times10^9$ M$_{\sun}$, traced by the CO(6-5) emission, together with the metal poor ([Fe/H]~$=-3.08 \pm 0.12$, [Si/H]~$=-2.86\pm0.14$) gas clump/filament which is possibly feeding its growth. The galaxy has an estimated star formation rate of $\approx100$ $M_{\sun}$ yr$^{-1}$, implying that it is a typical main sequence galaxy at these redshifts. The metal poor gas is detected through a damped Lyman-$\alpha$ absorber (DLA) observed at a spatial separation of 40 kpc and at the same redshift of Serenity-18, along the line of sight to the quasar SDSS J2310+1855 ($z_{\rm em}\simeq6.0025$). The chemical abundances measured for the damped Lyman-$\alpha$ system are in very good agreement with those measured for other DLAs discovered at similar redshifts, indicating an enrichment due to massive PopII stars. The galaxy/damped system that we discovered is a direct observational evidence of the assembly of a galaxy at the edge of the reionization epoch. | \label{sec:intro} Measuring the molecular gas content in early galaxies ($z\gtrsim 6$) is fundamental to reconstruct the cosmic star formation history, reionization, and enrichment. Molecular hydrogen is typically traced by emission from the carbon monoxide molecule, CO, because H$_2$ itself is generally not observable. While for a few galaxies at $z\approx6-8$ -- less than 1 billion years from the Big Bang -- we have sparse data on the stellar content (e.g. Bouwens et al. 2015; Jiang et al. 2016), the diffuse atomic gas (e.g. Carniani et al. 2017, 2018a) and the dust amount (e.g. Laporte et al. 2017; Hashimoto et al. 2018; Tamura et al. 2018), we completely lack information on their dense molecular component. At $z\sim6$, CO has been detected only in a few bright quasars tracing massive, highly star-bursting galaxies (Wang et al 2010, 2013; Gallerani et al. 2014; Venemans et al. 2017; Feruglio et al. 2018), which are rare objects that do not represent the bulk of the galaxy population in the early universe (e.g. Robertson et al. 2015). At lower redshift ($z\lesssim 4$), cold gas emission from galaxies has been recently detected in association with metal rich damped Lyman-$\alpha$ absorbers (DLAs), the strongest \ion{H}{1} absorptions observed in quasar spectra (characterized by column densities $N($\ion{H}{1}$) \ge 2 \times 10^{20}$ cm$^{-2}$). DLAs have the advantage that their detection is not biased by luminosity; on the other hand, the direct identification of their host galaxies has proven to be extremely challenging, in particular at high redshift (e.g. Fumagalli et al. 2017). Thanks to the lower luminosity of the quasar at sub-mm wavelengths, it was possible to observe with Atacama Large Millimeter/submillimeter Array (ALMA) four high-redshift ($z\sim2-4$), high-metallicity ($Z \sim0.1-1.0$ $Z_{\sun}$) DLAs revealing cold gas emission, either [\ion{C}{2}] or CO, in associated galaxies (Neeleman et al. 2017, 2018; Fynbo et al. 2018). The impact parameter between the galaxy detected in emission and the DLA varies between 18 and 117 kpc, implying that the DLA could arise in the circumgalactic medium of the target galaxy, but also that it could be associated with a fainter, undetected object. In this Letter we report the serendipitous discovery of a typical main sequence galaxy at $z\simeq 5.939$ (that we dubbed Serenity-18) detected in CO, associated with a metal poor DLA observed along the sightline to the quasar J231038.88+185519.7 ($z_{\rm em}=6.0025$, J2310 hereafter), using ALMA and XSHOOTER/Very Large Telescope (VLT) observations. Throughout this Letter we assume a standard flat $\Lambda$CDM cosmology, with $\Omega_\Lambda=0.7$ and $H_0=70$ \kms. | \subsection{DLA properties} On the basis of the large \ion{H}{1} column density measured for the DLA system, we assumed that no ionization corrections were needed and carried out the computation of the relative chemical abundances based on the column densities and the solar abundances from Asplund et al. (2009). The iron and silicon abundances, [Fe/H]~$=-3.08 \pm 0.12$ and [Si/H]~$=-2.86\pm0.14$, place this DLA absorber in the very metal poor regime as defined by Cooke et al. (2011). The absorption lines due to \ion{C}{2} $\lambda\,1334$ \AA\ and \ion{O}{1} $\lambda\,1302$ \AA\ could be saturated, thus we could derive only lower limits to the abundances of C and O. Previous studies of DLAs (e.g. Vladilo et al. 2011; Rafelski et al. 2012) have shown that below metallicities [Fe/H]~$\approx -2$ dust corrections are negligible; as a consequence, we are computing abundances assuming a dust-free gas. The studied absorber has abundances in very good agreement with those measured for the sample of $z\approx6$ DLAs by Becker et al. (2012). As already discussed in Becker et al. (2012), the chemical abundances observed for our system are also consistent with the 95 \% confidence interval of the abundances for the sample of very metal poor DLAs selected by Cooke et al. (2011) at $z\sim 2-4$\footnote{With the exclusion of the Carbon enhanced DLA along the line of sight to SDSS J0035-0918}. The observed abundance pattern is well explained by the predictions for PopII progenitors with $M \sim20$ $M_{\sun}$. The column density of \ion{H}{1} requires self-shielding of the gas from the cosmic ultraviolet (UV) background, which in turn implies that the absorbing gas density should be larger than $\simeq 0.1$ cm$^{-3}$ (Rahmati et al. 2013); this limit translates into a limit on the physical size of $\lesssim 4$ kpc. Together with the low metallicity, this suggests that we are seeing a gas filament/clump that has been recently forming from the intergalactic medium. \subsection{Serenity-18 properties} The CO-emitting galaxy is unresolved in our data. We estimate its upper limits size as $D \leq {\rm FWHM}_{\rm beam}=0.6$ arcsec, which corresponds to $\approx3.6$ kpc at the redshift of the galaxy. The inclination on the line of sight cannot be estimated either, because the source is unresolved. We derive the dynamical mass, modulus the inclination, by applying the relation, $M_{\rm dyn} \sin^2(i) = 1.16~ 10^5 \times (0.75\times $FWHM$_{\rm CO})^2 \times D$ (Wang et al. 2013, Feruglio et al. 2018), where $\rm FWHM_{\rm CO}=155$ \kms, and $D$ is the source size in kpc (diameter). From this relation we find, $M_{\rm dyn} \sin^2(i) \leq 5.6~10^9 ~ M_{\odot}$. In order to estimate the molecular gas mass from CO(6-5) we need to make some assumptions about the $S_{\rm CO}(6-5)/(1-0)$ ratio, and on the luminosity-to-mass conversion factor. We adopt in the following, $S_{\rm CO}(6-5)/(1-0)=r_{61}=20$, which is the average value measured for star forming galaxies from $z=0$ to 4 (Carilli \& Walter 2013). We note that an indication of a higher excitation at $z\sim6$, $r_{61}=70-150$, comes from recent hydrodynamical simulations by Vallini et al. (2018). Concerning the $\alpha_{\rm CO}$, the Milky Way value (i.e. 4.3 K \kms\ pc$^{-2}$ M$_{\sun}^{-1}$) is probably unphysical for high-$z$ galaxies. We therefore adopt a lower value based on Vallini et al. (2018), $\alpha_{\rm CO}=1.5$ K \kms\ pc$^{-2}$ M$_{\sun}^{-1}$. Based on these assumptions we infer a molecular gas mass of $M({\rm H}_2)= (5.4\pm0.5) \times 10^{9} \times (\alpha_{\rm CO}/1.5)~ M_{\sun}$. For different excitation properties of the gas (e.g. Vallini et al. 2018) the mass budget has to be decreased accordingly. For an inclination $i=90$ deg (edge-on) and $50$ deg, the molecular gas fraction in the galaxy would be $\mu= M({\rm H}_2)/M_{\rm dyn}\sim 0.9$ and 0.6, respectively. We note that the FWHM estimated by fitting a 1D Gaussian profile to the CO line, FWHM~$=155\pm30$ km s$^{-1}$, is similar to those estimated for [\ion{C}{2}] in typical galaxies at $z \approx 6 - 7$ (FWHM~$\approx150$ km s$^{-1}$, Maiolino et al. 2015; Knudsen et al. 2016; Pentericci et al. 2016; Bradac et al. 2017, Matthee et al. 2017; Carniani et al. 2018a,b). We cannot exclude, however, that part of the line may have been missed out of the bandpass (see Fig.~\ref{Fig:DLA_emiss}). In this case, our determinations of the gas mass and dynamical mass would be lower limits. From the calibration of Greve et al. (2014) between $L_{\rm FIR}$ and $L^{\prime}$ CO(6-5), we derive a far-infrared luminosity of $L_{\rm FIR} \approx 10^{12}~L_{\sun}$. This is a very rough estimate with an uncertainty of $\sim1$ dex due to the scatter of the correlation. We convert $L_{\rm FIR} $ to a star formation rate (SFR) using the Kennicutt et al. (1998) conversion factor corrected for a Kroupa (2011) initial mass function. We find $\rm SFR\approx115~ M_{\odot}/yr$. This value is consistent with the upper limit on the 150 $\mu$m continuum estimated from band 6 data. In fact, by assuming a typical SED of a star-forming galaxy with dust temperature in the range $\rm T_{dust}\sim 30-50~K$, we would expect a $3\sigma$ detection of the continuum for a $\rm SFR=200 ~M_{\odot}~yr^{-1}$ (see also Pallottini et al. 2017, Behrens et al. 2018). The properties we have derived for Serenity-18 allows us to state that this is the first detection of molecular gas emission from a typical main sequence galaxy at the end of the reionization epoch (see e.g. Santini et al. 2017). \subsection{The system Serenity-18 + DLA} The difference in redshift between the absorbing gas of the DLA and the emission line is of $\sim50$ \kms. Their angular separation is of $6.7$ arcsec, corresponding to an impact parameter of $\approx 40$ kpc at the DLA redshift, in a configuration similar to previous results for lower redshift DLAs (e.g. Neeleman et al. 2017, 2018). This confirms the association between the CO-emitting galaxy and the metal poor absorber. Furthermore, the relatively small redshift difference between the galaxy/DLA system and the quasar ($\Delta v \simeq -2746$ km s$^{-1}$ or $\approx 4$ proper Mpc) suggests that they could all be part of the same large-scale structure. State-of-the-art cosmological hydrodynamical simulations can help visualize the overall picture, which could give rise to our observations. We note that the luminosity of the CO(6-5) line, the inferred $M_{\rm dyn}$ and SFR of Serenity-18 are consistent with the predictions for the simulated $z\simeq6$ galaxy Alth{\ae}a (Pallottini et al. 2017; Vallini et al. 2018). In Fig.~\ref{Fig:sim}, we show the \ion{H}{1} column density and metallicity maps for the filament embedding Alth{\ae}a. In this context, the DLA can be identified with either one of the satellites of Alth{\ae}a, or with a gas condensation/filament surrounding the galaxy and possibly feeding it with fresh fuel for star formation. This picture is also in agreement with observations of clumpy galaxy assembly at $z\gtrsim6$ (Ouchi et al. 2010; Jiang et al. 2013; Bowler et al. 2017; Matthee et al. 2017; Carniani et al. 2018a, 2018b). The Serenity-18/DLA complex opens a new window in the study of typical galaxies in the early universe. It represents an ideal target for deeper, multiwavelength (UV/optical/NIR/mm) observations both in imaging and spectroscopy to understand the process of galaxy assembly. It suggests that also metal poor DLAs could be associated with CO-emitting galaxies: this could be true only at the highest redshifts, but it should be tested also at the lower redshifts where it is easier to select DLAs by metallicity. | 18 | 8 | 1808.01146 |
1808 | 1808.04887_arXiv.txt | Hydrogen-poor superluminous supernovae (SLSN-I) are a class of rare and energetic explosions discovered in untargeted transient surveys in the past decade\cite{qkk+11,gal12}. The progenitor stars and the physical mechanism behind their large radiated energies ($\sim10^{51}$ erg) are both debated, with one class of models primarily requiring a large rotational energy\cite{kb10,woo10}, while the other requires very massive progenitors to either convert kinetic energy into radiation via interaction with circumstellar material (CSM)\cite{wbh07,ci11,mbt+13,sbn16}, or engender a pair-instability explosion\cite{gmo+09,kbl+14}. Observing the structure of the CSM around SLSN-I offers a powerful test of some scenarios, though direct observations are scarce\cite{yqo+15,ylp+17}. Here, we present a series of spectroscopic observations of the SLSN-I iPTF16eh, which reveal both absorption and time- and frequency-variable emission in the Mg~II resonance doublet. We show that these observations are naturally explained as a resonance scattering light echo from a circumstellar shell. Modeling the evolution of the emission, we find a shell radius of 0.1~pc and velocity of 3300~km~s$^{-1}$, implying the shell was ejected three decades prior to the supernova explosion. These properties match theoretical predictions of pulsational pair-instability shell ejections, and imply the progenitor had a He core mass of $\sim 50-55~\Msun$, corresponding to an initial mass of $\sim 115~\Msun$. | 18 | 8 | 1808.04887 |
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1808 | 1808.05881_arXiv.txt | {BL Lac Objects (BL Lacs) and Flat Spectrum Radio Quasars (FSRQs) are radio-loud active galaxies (AGNs) whose jets are seen at a small viewing angle (blazars), while Misaligned Active Galactic Nuclei (MAGNs) are mainly radiogalaxies of type FRI or FRII and Steep Spectrum Radio Quasars (SSRQs), which show jets of radiation oriented away from the observer's line of sight. MAGNs are very numerous and well studied in the lower energies of the electromagnetic spectrum but are not commonly observed in the $\gamma$-ray energy range, because their inclination leads to the loss of relativistic boosting of the jet emission. The Large Area Telescope (LAT) on board the \emph{Fermi Gamma-ray Space Telescope} in the 100 MeV -- 300 GeV energy range detected only 18 MAGNs (15 radio galaxies and 3 SSRQs) compared to 1144 blazars. Studying MAGNs and their environment in the $\gamma$-ray sky is extremely interesting, because FRI and FRII radio galaxies are respectively considered the parent populations of BL Lacs and FSRQs, and these account for more than 50~$\%$ of the known $\gamma$-ray sources. The aim of this study is to hunt new $\gamma$-ray MAGN candidates among the remaining blazars of uncertain type and unassociated AGNs, using machine learning techniques and other physical constraints when strict classifications are not available. We found 10 new MAGN candidates associated with $\gamma$-ray sources. Their features are consistent with a source with a misaligned jet of radiation. This study reinforces the need for more systematic investigation of MAGNs in order to improve understanding of the radiation emission mechanisms and and the disparity of detection between more powerful and weaker $\gamma$-ray AGNs. } | Since 2008, the {\it Fermi Gamma-ray Space Telescope} has been observing active galactic nuclei (AGNs), providing an important window into the most extreme phenomenon of the universe: black holes. The Third {\it Fermi}-LAT Source Catalogue (3FGL) \citep{3fgl} reports $\gamma$-ray data collected in four years of operation (from 2008 August 11 to 2012 July 31) by the {\it Fermi} LAT (Large Area Telescope). The 3FGL Catalogue lists 3033 $\gamma$-ray sources within the 100 MeV--300 GeV energy range, of which 1745 are AGNs, mostly BL Lac objects (BL Lacs) and Flat Spectrum Radio Quasars (FSRQs), the two main kinds of blazars. AGNs with relativistic jets can be classified according to their jet orientation relative to the line of sight of the observer, as described in the unified model \citep{urry}. In the model, BL Lacs and FSRQs are active galaxies whose jets are seen at a small viewing angle. Another class of AGNs, Misaligned Active Galactic Nuclei (MAGNs), show a jet of radiation pointed at larger angles to the viewer. The radio-loud MAGNs are separated into two main classes \citep{urry,magn}, namely the edge-darkend Fanaroff-Riley I radio galaxies (FRI) and the edge-brightened Fanaroff-Riley II radio galaxies (FRII) \citep{fan} according to the distance of the brightest point from the central core. Steep Spectrum Radio Quasars (SSRQs), whose jet angles are smaller than those of radio galaxies but larger than those of blazars, are often also considered as MAGN \citep{ori}. We do not include SSRQs in the present analysis. There is no exact definition of what \emph{small viewing angle} means. In \citet{sb} a blazar is considered a source with a viewing angle $\theta_ {v}$$\leq$1/{$\Gamma$} where $\Gamma$ is the bulk Lorentz factor of the jet. BL Lacs and FSRQs are more easily detected with the {\it Fermi} LAT than MAGNs since the small inclination angle of the jets results in relativistic boosting of the non-thermal emission. MAGNs, which are detected in large numbers at radio and optical frequencies, are not so commonly observed in the $\gamma$-ray energy range, because their larger inclinations lead to the loss of this relativistic boosting. Only 11 FRI radio galaxies and 4 FRII radio galaxies (Table 1) are classified as MAGNs in the 3FGL Catalogue, compared to 1144 blazars. One interpretation of this result could be that {\it Fermi} cannot detect FRIIs because they are far away and therefore not bright enough to be detected. Other possible interpretations are studied by \citet{ele}, but the question remains open. More sources are necessary to fully understand this matter. MAGNs in the $\gamma$-ray sky remain extremely interesting, because FRI and FRII radio galaxies are respectively considered the parent populations of BL Lacs and FSRQs \citep{fr}, which account for more than 50~$\%$ of the known $\gamma$-ray sources. In principle, MAGN physics can significantly improve the knowledge of the relationship between the radiation emission mechanisms and the structure of the galaxies, as well as about the disparity of detection between more powerful and weaker $\gamma$-ray sources. The aim of this study is to hunt new MAGN candidates among the remaining blazars of uncertain type (BCUs) and unassociated AGNs (UCS$_{agn}$s) selected with the machine learning techniques used by \citet{bflap, pablo, zoo}. In studying MAGNs a great uncertainty remains about how to compute the jet angle. Recent studies by \citet{magn, giov, landt, paola} have introduced some methods for jet angle computing, but the results are still uncertain, and the topic remains open. Due to the uncertainty of the methods, we avoided detailed angle calculations and constrained our hunt to other physical and morphological features of MAGNs and when possible on optical spectroscopy. The paper is organized as follows: in Sect. 2 we provide a brief description of the {\it Fermi} LAT, and in Sect. 3 we describe the selection of MAGN candidates among the uncertain 3FGL objects. In Sect. 4 we present the multifrequency features of the most representative MAGN candidates and their optical spectra. In Sect. 5 we discuss the final results of this study. \begin{table} \begin{tabular}{lcccccr} \hline \hline \bf{3FGLname}&\bf{Association name} &\bf{FR type} &\bf{$L_{BLLac}$}\\ \hline 3FGL J0308.6$-$0408&NGC 1218& FRI & 0.80\\ 3FGL J0316.6+4119&IC310&FRI & 0.89\\ 3FGL J0319.8+4130&NGC1275&FRI & 0.70\\ 3FGL J0322.5$-$3712&NCG1316&FRI&0.90\\ 3FGL J0334.2+3915&4C+39.12&FRII & 0.26\\ 3FGL J0418.5+3013c&3C 111& FRII & 0.04\\ 3FGL J0519.2$-$4542&Pictor A&FRII & 0.14\\ 3FGL J0627.0$-$3529&PKS 0625-35&FRI & 0.96\\ 3FGL J0758.7+3747&NGC 2484& FRI& 0.83\\ 3FGL J1145.1+1935&3C 264&FRI&0.81\\ 3FGL J1230.9+1224&M87& FRI & 0.90\\ 3FGL J1325.4$-$4301& Cen A Core&FRI & 0.70\\ 3FGL J1346.6$-$6027& Cen B&FRI & 0.90\\ 3FGL J1442.6+5156& 3C 303& FRII & 0.98\\ 3FGL J1630.6+8232&NGC 6251& FRI & 0.98\\ \hline \end{tabular} \caption{Radio-galaxy MAGNs in the 3FGL {\it Fermi}-LAT Catalogue. } \end{table} | \label{sec:CON} Motivated by wishing to develop a process to expand the small sample of $\gamma$-ray MAGN, we have developed a search method based on machine learning combined with comparisons to parameters of the known $\gamma$-ray MAGN. The 10 identified candidates represent in many ways a first step toward that goal. Follow-up work is necessary to firmly establish these sources as MAGNs, including finding redshifts for half the candidates, using deeper radio, optical, and X-ray observations to reveal morphological details, and possibly undertaking Very Long Baseline Interferometry observations to examine the jets themselves. We also recognize that this search is necessarily incomplete. As already pointed out by \citet{3lac}, the $\gamma$-ray sources with unknown properties are generally fainter than the well-defined classes. The fainter sources offer less of the variability information required for the machine learning method, and so there may be $\gamma$-ray MAGN among the 529 3FGL sources eliminated in the first step of our method. The level of incompleteness is difficult to quantify with such a small number of sources. Nevertheless, the best candidates among our set are convincing as $\gamma$-ray MAGN and can be combined with the known 3FGL MAGN for population studies. Two of our candidates are also promising targets for TeV observations. The success of the method also indicates that it will be useful in searching for more $\gamma$-ray MAGN in the forthcoming 4FGL catalog, particularly when combined with the ongoing Very Large Array Sky Survey (VLASS)\footnote{\url{https://science.nrao.edu/science/surveys/vlass/}}. | 18 | 8 | 1808.05881 |
1808 | 1808.07556_arXiv.txt | { The $r_{01}$ and $r_{10}$ separation ratios are not independent so combing them into a single series $r_{010}$ is overfitting the data, this can lead to almost singular covariance matrices with very large condition numbers, and hence to spurious results when comparing models and observations. Since the $r_{02}$ ratios are strongly correlated with $r_{10}$ and $r_{01}$ ratios, they should be combined into a single series $r_{102}$ (or $r_{012}$), which are not overfitted, and models and observation compared using the covariance matrix $cov_{102}$ (or $cov_{012}$) of the combined set. I illustrate these points by comparing the revised Legacy Project data with my results on the 10 Kepler stars in common.} | Frequency separation ratios are widely used in asteroseismic model fitting, i.e. finding models whose oscillation properties match an observed set, as these ratios are almost independent of the outer layers of a star (Roxburgh and Vorontsov 2003, 2013). The ratios, constructed from frequencies $\nu_{n\ell}$ for angular degree $\ell=0,1,2$, are customarily defined as $$r_{01}(n)={(\nu_{n-1,0}-4\nu_{n-1,1}+6\nu_{n,0}-4\nu_{n,1}+\nu_{n+1,0})\over 8\, (\nu_{n,1}-\nu_{n-1,1})}~~~{\rm at}~~~\nu_{n.0}\eqno(1a)$$ $$r_{10}(n)=-{(\nu_{n-1,1}-4\nu_{n,0}+6\nu_{n,1}-4\nu_{n+1,0}+\nu_{n+1,1})\over 8\, (\nu_{n+1,0}-\nu_{n,0})}~~{\rm at}~\nu_{n,1}\eqno(1b)$$ $$r_{02}(n) = {\nu_{n,0} - \nu_{n-1,2}\over \nu_{n,1}-\nu_{n-1,1}} ~~~~{\rm at}~~~~\nu_{n,0}\eqno(1c)$$ Since the ratios for given $n$ have several frequencies in common (eg $\nu_{n,0}, \nu_{n,1}$), as do ratios of neigbouring $n$, they are strongly correlated. When comparing model and observed values this requires one to match models and observed values of the ratios using the covariance matrices of the ratios of the observed values. Care needs to be taken to ensure that one includes all the relevant correlations and that one does not overfit the data. | The ratios $r_{01}$, $r_{10}$ are not independent and in principle one set can be derived from the other by interpolation. From $N$ $\ell=0,1$ frequencies one can only derive $\sim N$ independent values of the phase shift differences $\delta_1-\delta_0$ (which are approximated by the ratios). But the $r_{010}$ sequence has $\sim 2N$ components. In this sense one is overfitting the data. Neighbouring elements of the covariance matrix can then be very strongly correlated leading to almost singular matrices with large condition numbers, and hence spurious results when comparing 2 sets of ratios. Only one of $r_{01}$, $r_{10}$ should be used in model fitting. \newpage Since the $r_{02}$, $r_{10}$ and $r_{01}$ ratios are correlated they should be combined into single sequence $r_{012}$ or $r_{102}$ when comparing 2 sets of ratios. These sequences are not overfitted since from $N$ $\ell=0,1,2$ frequencies subtraction gives $\sim N$ values of both $\delta_1-\delta_0$ and $\delta_2-\delta_0$, which are approximated by the ratios $r_{10}, r_{02}$. The $r_{102}$ covariance matrices have reasonable condition numbers. \begin{table} [t] \setlength{\tabcolsep}{8pt} \caption {Fits of LegacyN ratios and covariances to Roxburgh's ratios } \vskip -6pt \small \centering \begin{tabular}{l c c c c c c c c c c r c c c } \hline\hline \noalign{\smallskip} KIC no&$$& $r_{10}$ & $r_{02}$ &$r_{102}$ & $r_{012}$\\ [0.5ex] \hline \noalign{\smallskip} 3427720 & $\chi^2_{cov}$ & 0.427 & 0.772 & 0.582 & 0.580 \\[0.5ex] 6106415 & $\chi^2_{cov}$ & 1.198 & 2.507 & 2.206 & 2.188 \\[0.5ex] 6116048 & $\chi^2_{cov}$ & 0.657 & 0.597 & 0.562 & 0.567 \\[0.5ex] 6225718 & $\chi^2_{cov}$ & 0.830 & 1.397 & 1.060 & 1.042 \\[0.5ex] 6603624 & $\chi^2_{cov}$ & 0.138 & 1.926 & 1.162 & 1.226 \\[0.5ex] 8379927 & $\chi^2_{cov}$ & 0.368 & 2.958 & 1.576 & 1.544 \\[0.5ex] 8760414 & $\chi^2_{cov}$ & 0.661 & 2.831 & 1.742 & 1.550 \\[0.5ex] 9098294 & $\chi^2_{cov}$ & 0.481 & 0.244 & 0.430 & 0.402 \\[0.5ex] 12069424 & $\chi^2_{cov}$ & 1.059 & 0.857 & 0.919 & 0.947 \\[0.5ex] 12069449 & $\chi^2_{cov}$ & 1.577 & 1.022 & 1.360 & 1.225 \\[0.5ex] \hline \end{tabular} \end{table} \begin{table} [t] \setlength{\tabcolsep}{8pt} \caption {Fits of Roxburgh's ratios and covariances to LegacyN ratios } \vskip -6pt \small \centering \begin{tabular}{l c c c c c c c c c c r c c c } \hline\hline \noalign{\smallskip} KIC no&$$& $r_{10}$ & $r_{02}$ &$r_{102}$ & $r_{012}$\\ [0.5ex] \hline \noalign{\smallskip} 3427720 & $\chi^2_{cov}$ & 0.489 & 0.683 & 0.544 & 0.536 \\[0.5ex] 6106415 & $\chi^2_{cov}$ & 1.270 & 2.762 & 2.259 & 2.361 \\[0.5ex] 6116048 & $\chi^2_{cov}$ & 0.950 & 0.735 & 0.783 & 0.793 \\[0.5ex] 6225718 & $\chi^2_{cov}$ & 0.663 & 1.045 & 0.859 & 0.924 \\[0.5ex] 6603624 & $\chi^2_{cov}$ & 0.162 & 4.139 & 2.591 & 2.798 \\[0.5ex] 8379927 & $\chi^2_{cov}$ & 0.441 & 1.069 & 0.701 & 0.646 \\[0.5ex] 8760414 & $\chi^2_{cov}$ & 1.938 & 10.322 & 5.557 & 5.523 \\[0.5ex] 9098294 & $\chi^2_{cov}$ & 0.564 & 0.268 & 0.479 & 0.425 \\[0.5ex] 10963065 & $\chi^2_{cov}$ & 1.524 & 1.821 & 1.535 & 1.496 \\[0.5ex] 12069424 & $\chi^2_{cov}$ & 2.607 & 6.608 & 4.697 & 4.790 \\[0.5ex] 12069449 & $\chi^2_{cov}$ & 1.935 & 6.090 & 4.883 & 3.965 \\[0.5ex] \hline \end{tabular} \end{table} | 18 | 8 | 1808.07556 |
1808 | 1808.01683.txt | We use joint observations by the \emph{Neil Gehrels} \Swift X-ray Telescope (XRT) and the \Fermi Large Area Telescope (LAT) of gamma-ray burst (GRB) afterglows to investigate the nature of the long-lived high-energy emission observed by \Fermi LAT. Joint broadband spectral modeling of XRT and LAT data reveal that LAT nondetections of bright X-ray afterglows are consistent with a cooling break in the inferred electron synchrotron spectrum below the LAT and/or XRT energy ranges. Such a break is sufficient to suppress the high-energy emission so as to be below the LAT detection threshold. By contrast, LAT-detected bursts are best fit by a synchrotron spectrum with a cooling break that lies either between or above the XRT and LAT energy ranges. We speculate that the primary difference between GRBs with LAT afterglow detections and the non-detected population may be in the type of circumstellar environment in which these bursts occur, with late-time LAT detections preferentially selecting GRBs that occur in low wind-like circumburst density profiles. Furthermore, we find no evidence of high-energy emission in the LAT-detected population significantly in excess of the flux expected from the electron synchrotron spectrum fit to the observed X-ray emission. The lack of excess emission at high energies could be due to a shocked external medium in which the energy density in the magnetic field is stronger than or comparable to that of the relativistic electrons behind the shock, precluding the production of a dominant synchrotron self-Compton (SSC) component in the LAT energy range. Alternatively, the peak of the SSC emission could be beyond the 0.1--100 GeV energy range considered for this analysis. | Joint observations by NASA's \emph{Neil Gehrels} \Swift and \Fermi missions have led to a unique opportunity to study the broadband properties of gamma-ray bursts (GRBs) over an unprecedentedly broad energy range. The two missions have the combined capability of probing the emission from GRBs over 11 decades in energy, ranging from optical ($\sim$2 eV) to high-energy gamma rays ($>300$ GeV). After more than 7 yrs of simultaneous operations, \Swift and \Fermi have detected thousands of GRBs, with over 100 of these bursts detected at energies greater than 30 MeV by the \Fermi Large Area Telescope (LAT) \citep{Vianello2015}\footnote{https://fermi.gsfc.nasa.gov/ssc/observations/types/grbs/lat\_grbs/}. %The high-energy emission observed by the LAT is typically temporally extended, lasting longer than the emission observed at keV energies by the \Swift Burst Alert Telescope (BAT) and the \Fermi Gamma-ray Burst Monitor (GBM). There also appears to be a consistent delay in the onset of the LAT-detected emission with respect to the emission observed at lower energies. The delayed onset and long-lived nature of the LAT detected emission suggests that the afterglow components commonly observed in X-ray, optical, and radio wavelengths may also produce a significant amount of gamma-ray emission \citep{KumarBarniolDuran2009, Ghisellini2010, DePasquale2010}. In this scenario, the late-time emission detected by the LAT would be due to the high-energy extension of the electron synchrotron spectrum produced by the external forward shock associated with the GRB blast wave moving into the circumstellar environment. Broadband fits to the simultaneous multiwavelength observations of GRB 110731A \citep{GRB110731A} and GRB 130427A \citep{GRB130427A_LAT} show similar late-time spectral and temporal behavior, supporting such an external shock interpretation. Likewise, a stacking analysis of the LAT data of \Swift localized bursts that were not detected above 50 MeV has shown evidence for a subthreshold signal that correlates with the X-ray brightness of their afterglow emission as measured by Swift-XRT, lending further support for an afterglow origin to the observed gamma-ray emission. %Observations by NASA \emph{Fermi Gamma-ray Space Telescope} have led to an unique opportunity to study the broadband properties of gamma-ray bursts (GRBs) over an unprecedented energy range. With its two primary instruments, the Gamma-ray Burst Monitor and the Large Area Telescope, the spacecraft has the combined capability of probing the emission from GRBs over 8 decades in energy, ranging from $\sim8$ keV to $>300$ GeV. In 7 years or triggered operation, the \Fermi GBM has detected thousands of GRBs, with over 100 of these bursts detected at energies greater than 50 MeV by \Fermi LAT. The properties of the high-energy emission observed by the LAT can differ considerably from the emission detected at keV and MeV energies by other instruments. While some bursts show evidence for emission in coincidence with activity at keV and MeV energies as observed by the \Swift Burst Alert Telescope (BAT) and \Fermi Gamma-ray Burst Monitor (GBM) \citep{Ackermann2010}, others also exhibit high-energy emission that is temporally extended, lasting longer than the emission observed at lower energies \citep{GRB110731A, GRB130427A_LAT}. There also appears in some cases to be a delay in the onset of the LAT-detected emission with respect to the emission observed at lower energies \citep{Abdo2009a, Abdo2009b, Ackermann2013}. The delayed onset and long-lived component of the LAT-detected emission suggest that GRB afterglows commonly observed in X-ray, optical, and radio wavelengths may also produce significant gamma-ray emission \citep{KumarBarniolDuran2009, Razzaque2010a, Ghisellini2010, DePasquale2010}. In this interpretation, the coincident emission detected by the LAT is thought to be an extension of the prompt emission spectrum commonly attributed to shocks internal to the relativistic outflow \citep{Ackermann2010, Maxham2011, Zhang2011, Yassine2017}, while the late-time emission is due to the high-energy extension of the electron synchrotron spectrum produced by the external forward shock associated with the GRB blast wave moving into the circumstellar environment. %The properties of the high-energy emission observed by the LAT differ considerably from the emission detected at keV and MeV energies by other instruments. The high-energy emission is typically temporally extended, lasting longer than the emission observed at keV energies by both the \Swift Burst Alert Telescope (BAT) and \Fermi Gamma-Ray Burst Monitor (GBM). There also appears to be a consistent delay in the onset of the LAT-detected emission with respect to the emission observed at lower energies \citep{Ackermann2013}. The delayed onset and long-lived nature of the LAT-detected emission suggest that the afterglow components commonly observed in X-ray, optical, and radio wavelengths may also produce significant gamma-ray emission \citep{KumarBarniolDuran2009, Razzaque2010a, Ghisellini2010, DePasquale2010}. In this scenario, the late-time emission detected by the LAT is due to the high-energy extension of the electron synchrotron spectrum produced by the external forward shock associated with the GRB blast wave moving into the circumstellar environment. Broadband fits to the simultaneous multiwavelength observations of GRB 110731A \citep{GRB110731A} and GRB 130427A \citep{GRB130427A_LAT} show similar late-time spectral and temporal behavior, supporting such an external shock interpretation. Likewise, a stacking analysis of the LAT data of {\it Swift}-localized bursts that were not detected above 40 MeV has shown evidence for subthreshold emission on timescales that far exceed the typical duration of the prompt emission at keV energies \citep{Beniamini2011, LATStackingAnalysis}. Furthermore, the strength of this high-energy subthreshold emission correlates directly with the X-ray brightness of the burst's afterglow emission, as measured by the \Swift X-ray Telescope (XRT). Despite the growing evidence for an external shock origin of the long-lived high-energy emission observed by the LAT, the fact remains that only $\sim8\%$ of the bursts detected at keV energies within the LAT field of view (FoV) have been detected above 40 MeV \citep{Ackermann2013}. Therefore, although the signature of the afterglow emission at X-ray wavelengths is largely ubiquitous in GRBs observed by the XRT, the high-energy component is observed in only a small subset of these bursts. This has led to speculation that LAT-detected bursts may represent a unique population of GRBs, either probing a particular type of environment \citep{Racusin2011, Beloborodov2014}, the result of a unique set of afterglow conditions \citep{Ghisellini2010}, or the result of progenitors that produce a rare class of hyperenergetic GRBs \citep{Cenko2011}. %In this paper we systematically investigate the the long-lived high-energy emission associated with GRBs through the use of broadband data collected by \Swift XRT and \Fermi LAT. %By examining joint \Swift and \Fermi observations for over 380 Swift detected GRBs, we can characterize the high-energy spectral energy distribution of the LAT detected and non-detected populations and determine if the LAT-detected bursts differ significantly in their afterglow properties from the general GRB population. %By examining joint \Swift and \Fermi observations for over 380 Swift detected GRBs, we can characterize the high-energy spectral energy distribution of the LAT detected and non-detected populations and determine if the LAT-detected bursts differ significantly in their afterglow properties from the general GRB population. %In this paper we systematically investigate the long-lived high-energy emission associated with GRBs through the use of broadband data collected by \Swift XRT and \Fermi LAT. By examining joint \Swift and \Fermi observations for over 380 Swift detected GRBs, we can characterize the broadband spectral energy distribution (SED) of the LAT detected and non-detected afterglow populations. From these SED models, we can determine if the relative sensitivity of the XRT and LAT is sufficient to account for the majority of LAT nondetections of bright X-ray afterglows, or whether the LAT-detected bursts differ significantly in their afterglow properties from the general GRB population. %In this paper we attempt to address this question regarding the nature of LAT detected GRBs through the use of broadband data of bursts detected by both the \Swift and \Fermi spacecraft. By examining joint \Swift XRT and \Fermi LAT observations for over 380 GRBs, we can model the broadband spectral energy distribution (SED) of the afterglow emission associated with LAT detected and non-detected GRBs. We can then use these SEDs to determine if the relative sensitivity of the XRT and LAT is sufficient to account for the majority of LAT nondetections, or whether the LAT-detected bursts differ significantly in their afterglow properties from the general GRB population. In this paper we attempt to address the conditions that are required to produce the late time high-energy emission detected by the LAT through the use of broadband data collected by both \Swift and \emph{Fermi}. By examining joint XRT and LAT observations of 386 GRBs from 2008 August 4 to 2014 March 23, we can model the broadband spectra of the afterglow emission associated with LAT-detected and non-detected GRBs. This allows us to determine whether the relative sensitivities of the XRT and LAT are sufficient to account for the majority of LAT nondetections, or whether the LAT-detected bursts differ significantly in their afterglow properties from the general GRB population. A subset of these bursts is also subjected to detailed broadband spectral fitting of the simultaneous XRT and LAT data. From these spectral fits, we can determine whether the XRT and LAT data are consistent with being drawn from the same power-law segment (PLS) of an electron synchrotron spectrum, or if a break or suppression of the high-energy emission is required to explain the LAT non-detection. This analysis also allows us to place constraints on the existence of spectral components at high energies that are in excess of that predicted by the electron synchrotron model, such as external inverse Compton (EIC) \citep{Fan2006, He2012, Beloborodov2014A} and synchrotron self-Compton (SSC) \citep{Dermer2000, Zhang2001, Sari2001, Wang2013} contributions. %We find that the relative sensitivity between the XRT and LAT is sufficient to account for the majority of LAT nondetections of \Swift detected GRBs. For the small subsample of LAT nondetections of bright X-ray afterglows for which their XRT derived spectra predict emission in the LAT energy range, we find that a break in the SED, consistent with a cooling break in the predicted electron synchrotron spectrum, near or below the XRT energy range, is sufficient to suppress the high-energy emission so as to be below the LAT detection threshold. An analysis of the broadband data of LAT detected bursts, on the other hand, reveals that these bursts tend to be drawn from the brightest and hardest afterglows observed by the XRT. Moreover, all of the LAT detections we considered can be accommodated with an electron synchrotron spectrum where the cooling break either lies between or above the XRT and LAT energy ranges. We discuss the implications these results have on the type of circumstellar environment that would be required to produce the late time high-energy emission detected by the LAT. Finally, we find no evidence of high-energy emission in excess of what is predicted from an electron synchrotron spectrum, allowing us to place constraints on the ubiquity of inverse Compton (IC) components at late times. The paper is structured as follows: in \S\ref{sec:InstrumentOverview}, we review the characteristics of the \Fermi LAT and \Swift XRT instruments. In \S\ref{sec:SampleDefinition}, we define the GRB samples considered in this work and outline the analysis performed in \S\ref{sec:Analysis}. We present the results in \S\ref{sec:Results}, and discuss the implications of our results in \S\ref{sec:Discussion}. Unless specified otherwise, all temporal and spectral indices are defined as $F_\nu\propto E^{-\beta}t^{-\alpha}$, where $\beta=\Gamma-1$, with $\Gamma$ is the photon index. | \label{sec:Conclusions} We have used joint observations by the \Swift XRT and the \Fermi LAT of GRB afterglows to investigate the nature of long-lived, high-energy emission observed by \Fermi LAT. By extrapolating the XRT derived spectra of \emph{Swift}-detected GRBs, we compared the expected flux in the 0.1 to 100 GeV energy range to the LAT upper limits for the periods in which the burst position was within the LAT FoV. We found that only a small subset of bursts exhibit afterglow emission that could exceed the LAT detection threshold when extrapolated to the 0.1 to 100 GeV energy range. Bursts that do result in late-time LAT detections are almost exclusively drawn from afterglows that exhibit emission brighter than $F_{\rm XRT} \gtrsim 10^{-10}$ erg cm$^{-2}$ s$^{-1}$ and harder than $\Gamma_{\rm XRT} \lesssim 2$. Joint broadband spectral fits of XRT and LAT data reveal that a majority of LAT nondetections of relatively bright X-ray afterglows can be explained by an afterglow spectrum with a slightly softer photon index when constrained by both the XRT and LAT data, compared to the photon index derived by fits to the XRT data alone. The remaining LAT nondetections are consistent with a cooling break in the predicted electron synchrotron spectrum between the XRT and LAT energy ranges. Such a break is sufficient to suppress the high-energy emission below the LAT detection threshold. On the other hand, the broadband spectra of LAT-detected bursts are best modeled by spectral components that indicate that the cooling break in the synchrotron spectrum lies either between or above the XRT and LAT energy ranges. Since the value and time evolution of the cooling frequency in an electron synchrotron spectrum is strongly dependent on the density profile of the circumstellar medium, we speculate that the primary difference between bursts with afterglow detections by the LAT and the non-detected population may be the type of circumstellar environment. Late-time LAT detections may be preferentially selecting GRBs that occur in low-density wind-like circumburst environments for which the synchrotron cooling break begins near the X-ray regime and does not evolve to lower energies, resulting in an afterglow spectrum above the X-ray regime that remains spectrally hard for longer periods of time, enhancing the detectability of the afterglow in the LAT energy range. We find no evidence of high-energy emission significantly in excess of the flux expected from the spectrum predicted by the electron synchrotron model. In addition, joint spectral fits of contemporaneous XRT and LAT observations of an episode of energetic X-ray flaring in GRB~100728A and a significant X-ray plateau in GRB~110213A find that the XRT and LAT data are consistent with a single spectral component. The lack of excess emission at high energies points to two possibilities: 1) a shocked external medium in which the energy density in the magnetic field is elevated or comparable to that of the relativistic electrons behind the shock, precluding the production of a dominant SSC component in the LAT energy range at late times, or 2) the peak of the SSC emission is beyond the 0.1 to 100 GeV energy range we considered. \medskip \noindent The \textit{Fermi} LAT Collaboration acknowledges generous ongoing support from a number of agencies and institutes that have supported both the development and the operation of the LAT as well as scientific data analysis. These include the National Aeronautics and Space Administration and the Department of Energy in the United States, the Commissariat \`a l'Energie Atomique and the Centre National de la Recherche Scientifique / Institut National de Physique Nucl\'eaire et de Physique des Particules in France, the Agenzia Spaziale Italiana and the Istituto Nazionale di Fisica Nucleare in Italy, the Ministry of Education, Culture, Sports, Science and Technology (MEXT), High Energy Accelerator Research Organization (KEK) and Japan Aerospace Exploration Agency (JAXA) in Japan, and the K.~A.~Wallenberg Foundation, the Swedish Research Council and the Swedish National Space Board in Sweden. Additional support for science analysis during the operations phase is gratefully acknowledged from the Istituto Nazionale di Astrofisica in Italy and the Centre National d'\'Etudes Spatiales in France. This work performed in part under DOE Contract DE-AC02-76SF00515. \clearpage % ============= Tables ============= %\include{Table1} % ============= References ============= | 18 | 8 | 1808.01683 |
1808 | 1808.00971_arXiv.txt | In a $\Lambda$CDM cosmology, galaxy formation is a globally inefficient process: it is often the case that far fewer baryons are observed in galaxy disks than expected from the cosmic baryon fraction. The location of these ``missing baryons'' is unclear. By fitting halo profiles to the rotation curves of galaxies in the SPARC data set, we measure the ``missing baryon'' mass for individual late-type systems. Assuming that haloes initially accrete the cosmological baryon fraction, we show that the maximum energy available from supernovae is typically not enough to completely eject these ``missing baryons'' from a halo, but it is often sufficient to heat them to the virial temperature. The energy available from supernovae has the same scaling with galaxy mass as the energy needed to heat or eject the ``missing baryons'', indicating that the coupling efficiency of the feedback to the ISM may be constant with galaxy virial mass. We further find that the energy available from supernova feedback is always enough to convert a primordial cusp into a core and has magnitude consistent with what is required to heat the ``missing baryons'' to the virial temperature. Taking a census of the baryon content of galaxies with ${\rm 10^9<M_{vir}/M_{\odot}<10^{12}}$ reveals that $\sim86\%$ of baryons are likely to be in a hot phase surrounding the galaxies and possibly observable in the X-ray, $\sim7\%$ are in the form of cold gas, and $\sim7\%$ are in stars. | In the early Universe, it is predicted that dark matter and baryons are well mixed \citep{Spergel2003}. As dark matter haloes form by gravitational collapse, the baryons cool, dissipate energy, and fall to the centres of the haloes \citep{White1978,Fall1980}. Assuming a simple model where no other processes are at play, the total mass of a galaxy should then be composed of $\sim15\%$ baryons and $\sim85\%$ dark matter \citep{Planck2015}. However, only in the most massive galaxy clusters does the observed baryon fraction begin to approach this value \citep{Giodini2009,Gonzalez2013}, while at dwarf galaxy masses, only $\sim1\%$ of the expected baryon content is actually detected \citep{McGaugh2010,Papastergis2013,Bradford2015}. A simple observational census of the baryonic content of the Universe therefore reveals that most baryons have not been identified. Many of these ``missing baryons" may reside in the inter-galactic medium (IGM), although this is unlikely to account for the full expected quantity \citep{Fukugita1998,Danforth2005,Shull2012}. Alternatively, these ``missing baryons" may exist in a phase that is difficult to detect \citep{Bregman2015}. Accretion onto dark matter haloes can shock heat the gas up to the virial temperature (${\rm T_{vir}}$), which would make the gas emit primarily in the X-ray \citep{White1991,Cen1999}. Many hydrodynamic simulations show that galaxies, especially at high redshift, are primarily fed by cold flows which penetrate deep into the centres of the haloes \citep{Keres2005}. However, this may not be the case at low redshift. Strong feedback processes resulting from galaxy formation can reheat much of this gas into a hot phase, placing the primary emission in the X-ray \citep{Mathews1971,McKee1977,Cen1999}. Hot gaseous haloes have been detected around a number of galaxies \citep{Anderson2011,Anderson2013,Miller2015,Bregman2018}, although for massive spirals, this may not completely account for the entire budget of ``missing baryons" \citep{Li2018,Bregman2018}. Hot baryons may also be detected by the thermal Sunyaev-Zel'dovitch (kSZ) effect \citep{tSZ1, tSZ2}. As observations improve more of the ``missing baryons" are being found, but it is unclear whether they are sufficient to make up the baryon content of the Universe. This inefficiency by which galaxies either obtain or retain their baryons translates directly into the stellar content of galaxies. Abundance matching techniques predict that even for the most efficient star-forming galaxies, the fraction of stars present remains far below the cosmic baryon fraction at all masses \citep{Moster2013, Behroozi2013, Moster2017}. Explanations for the low fraction of stars come from feedback in three regimes. At the very lowest halo masses, (${\rm M_{halo}\lesssim10^9\ M_{\odot}}$), reionization may have completely prevented galaxies from forming by limiting their ability to accrete gas \citep{Babul1992,Efstathiou1992,Gnedin2000,Okamoto2008}. At slightly higher masses, supernova (SN) feedback can limit star formation by either ejecting gas from the halo or keeping it at a temperature ${\rm T \sim T_{vir}}$ \citep{Dekel1986}. Finally, at the highest mass systems, feedback from AGN may be the dominant mechanism which regulates galaxies \citep{Silk1998,Bower2006}. Feedback processes can impact galaxies in many different ways. For example, the slope of the predicted black hole mass--stellar velocity dispersion relation depends on whether AGN outflows are ``momentum-driven" or ``energy-driven" \citep{King2003,Costa2014}. For stellar feedback, the impulsiveness of the energy injection can determine whether or not a dark matter cusp can be converted into a core \citep{Navarro1996b,Pontzen2012}. The way in which the feedback is modelled has drastic effects on observable quantities such as the stellar mass function and the galaxy distribution along the Hubble sequence \citep[e.g.][]{Schaye2015,Vogelsberger2014,Dubois2014,Hopkins2017}, bulge formation \citep[e.g.][]{Hopkins2012}, central black hole mass \citep[e.g.][]{Curtis2016,DiCintio2017}, gas fractions and X-ray luminosities \citep[e.g.][]{Puchwein2008}. Understanding the energy scales involved and how the feedback energy couples to the local medium is key to developing a complete model for galaxy formation. Historically, one of the most well studied effects of feedback is how the density profile of the dark matter halo responds to large scale inflows or outflows of baryons \citep{Navarro1996,Pontzen2012,Blumenthal1986,Gnedin2004,Katz2014,Martizzi2013,Read2005,Read2016}. \cite{Penarrubia2012} calculated the range of energies required to transform a cuspy density profile to one with a core and compared it to the energy budget of SN. They identified that Milky Way dSphs require $10^{53-55}$~ergs of energy in order to form cores with sizes comparable to the luminous size of the galaxies. Using similar analytic reasoning, we focus here on how stellar feedback impacts the slope of the ${\rm M_*-M_{halo}}$ relation for late-type galaxies (spirals and dIrr) as well as dark matter density profiles. We focus on two halo profiles: NFW \citep{Navarro1996}, which is a prediction from cosmological DM-only simulations, and DC14 \citep{DC2014,DiCintio2014b}, which is a fit to the halos produced in the MaGICC simulations \citep{Brook2012,Stinson} that include the effects of galaxy formation. This model includes a dependence of halo shape on ${\rm M_*/M_{halo}}$ such that the dark matter profiles of low and high mass galaxies are cuspy at the centre, while intermediate mass galaxies exhibit cores. Although this model is successful at accounting for the kinematics of a range of galaxy types \citep{DC2014,Katz2016} it is important to bear in mind that the formation of cores is a contentious issue among the current-generation of cosmological hydrodynamic simulations. For example, although other high-resolution simulations produce cores of roughly the DC14 type \citep[e.g.][]{Chan, Read}, others do not produce cores at all \citep[e.g.][]{Apostle}. The exact mechanism of core formation has yet to be determined and there are indeed degeneracies in subgrid implementations of star formation and feedback that may lead to the same results. Nevertheless, since the DC14 model fits the properties of observed rotation curves well, we aim to use the energy scales of real galaxies as a further test of whether SN can provide enough energy to both regulate star formation and produce cores in real galaxies. We begin by presenting observational results on the fraction of cold gas and stars of $\sim150$ late-type galaxies from the SPARC data set \citep{Lelli2016} (Sec.~\ref{obscon}). We then review the theoretical motivation for how SN feedback regulates galaxy formation and show that in two simple models it naturally leads to a logarithmic slope of $5/3$ for the ${\rm M_*-M_\text{vir}}$ relation (Sec.~\ref{theory}). These models are then tested against the SPARC data where this slope can be independently measured: we then check whether this energy budget is sufficient to generate a core in low mass galaxies or reverse adiabatic contraction in higher mass systems (Sec.~\ref{comparison}). Finally, in Sec.~\ref{discussion} we use these results to take a cosmic census of the distribution of baryons in late-type galaxies with $10^9<{\rm M_{vir}}<10^{12}$. Throughout this work, we assume a WMAP3 cosmology with $H_0 = 73{\rm km~s^{-1}~Mpc^{-1}}$, $\Omega_{\rm m} = 0.24$, $\Omega_{\Lambda}= 0.76$, $\Omega_{\rm b} = 0.04$ and $\sigma_8 = 0.76$ \citep{Spergel2007}, and we define the virial radius to be the radius which contains a mean density equal to $93.6\:\rho_{\rm crit}$. | We have calculated the fraction of ``missing baryons" in 147 galaxies from the SPARC database, using empirically estimated halo masses from rotation curve fits for two different halo models: NFW and DC14 (see \citealt{Katz2016}). Our main results are as follows: \begin{itemize} \item{} We confirm that galaxy formation is a globally inefficient process. The observed fraction of baryons present in the cold disk is in general much smaller than the cosmic baryon fraction. This fraction scales weakly with the virial mass of the galaxy such that $f_d\propto{\rm M_{vir}^{0.16}}$, although there is significant scatter about this relation. Likewise, the fraction of baryonic mass comprised of stars scales strongly with mass; the most massive galaxies in our sample are completely dominated by stars. At ${\rm M_{vir}\sim10^{11.1}}$, roughly 50\% of the observed baryonic content is in cold gas and the other 50\% is in stars. \item{} When comparing the energy required to heat the ``missing baryons" to the virial temperature versus eject them from the halo, we find that far less energy is required for the former. For SN feedback to regulate the heating, $\lesssim10\%$ of the total SN energy available needs to couple to the gas (dependent on the stellar IMF). We find that a constant coupling efficiency is sufficient to explain the ``missing baryons" for late-type spiral galaxies in the mass range $10^9<{\rm M_{vir}/M_{\odot}}<10^{12}$ when haloes are modelled with a DC14 profile. This then sets the slope of the ${\rm M_*/M_{vir}}$ relation to be slightly less than $5/3$ due to the weak scaling of $f_d$ with ${\rm M_{vir}}$. However, there is typically not enough SN energy to eject all of the missing baryons from the halo. \item{} When comparing the energy required to restructure the halo from a primordial NFW profile to the DC14 form, we find that there is always more energy available from SN feedback than required for the transformation. We find that there is no mass dependence in the SN coupling efficiency required to create cores and that the magnitude of this efficiency is very similar to what is required to heat the ``missing baryons". In contrast, there is a mass dependency in the efficiency required to erase the halo transformation for systems that exhibit contraction. This may indicate that halo restructuring is dependent on quantities other than stellar mass such as halo mass, formation redshift, or star formation history. \item{} We take a cosmic census of the baryons that should be associated with late-type galaxies, finding that $\sim86\%$ of the total baryonic content is likely to be in a hot halo surrounding the galaxy, and the remaining $14\%$ split equally between stars and gas. In order for our predictions to be confirmed, future X-ray or kinetic Sunyaev-Zeldovich observations will have to accurately measure the masses in hot gas surrounding these local galaxies. \end{itemize} | 18 | 8 | 1808.00971 |
1808 | 1808.02659_arXiv.txt | A period of early matter domination can give rise to the correct dark matter abundance for a broad range of dark matter annihilation rate $\langle \sigma_{\rm ann} v \rangle_{\rm f}$. Here, we examine this scenario for situations where $\langle \sigma_{\rm ann} v \rangle_{\rm f}$ is below the nominal value for thermal dark matter $3 \times 10^{-26}$ cm$^3$ s$^{-1}$ as possibly indicated by some recent experiments. We show that obtaining the correct relic abundance sets a lower bound on the duration of early matter domination era in this case. On the other hand, provided that the post-inflationary universe has an equation of state characterized by $w \leq 1/3$, the requirement that the scalar spectral index $n_s$ be within the observationally allowed range limits the duration of this epoch from above. By combining these considerations, we show that the current and future cosmic microwave background experiments can tightly constrain the parameter space for this scenario. In particular, models of inflation with a tensor-to-scalar ratio below ${\cal O}(0.01)$ may disfavor non-thermal supersymmetric dark matter from a modulus-driven early matter domination epoch. | Despite various lines of evidence for the existence of dark matter (DM) in the universe~\cite{BHS}, its identity remains as a major problem at the interface of cosmology and particle physics. Weakly interacting massive particles (WIMPs) are promising candidates for DM and are the main focus of direct, indirect, and collider searches that are currently underway to discover DM. A nice mechanism for obtaining the correct abundance for WIMP DM is the "WIMP miracle", which assumes that the universe was in a radiation-dominated (RD) phase at temperatures about the DM mass $m_\chi$. The DM relic abundance in this picture is set when the annihilation rate of DM particles drops below the Hubble expansion rate, called "thermal freeze-out", and matches the observed value if the annihilation rate takes the nominal value $\langle \sigma_{\rm ann} v \rangle_{\rm f} = 3 \times 10^{-26}$ cm$^3$ s$^{-1}$. However, the WIMP miracle has come under increasing scrutiny by the recent experimental data. For example, Fermi-LAT's results from observations of dwarf spheroidal galaxies~\cite{fermi1} and newly discovered Milky Way satellites~\cite{fermi2} place an upper bound on the DM annihilation rate that is below the nominal value required for the WIMP miracle for a range of DM masses. A recent analysis~\cite{Beacom} shows that in models where DM annihilation is dominated by $S$-wave processes, thermal DM with a mass below 20 GeV is ruled out in a model-independent way, while for certain annihilation channels this will be the case for masses up to 100 GeV. This implies that thermal freeze-out in a RD universe would lead to overproduction of DM within the corresponding mass range (unless there is $P$-wave annihilation or co-annihilation, in which cases the WIMP miracle condition and the indirect detection limits may be both satisfied). However, the situation can change in a non-standard thermal history where the universe is not RD at the time of freeze-out~\cite{KT}. In particular, it is known that an epoch of early matter domination (EMD) that ends before the onset of big bang nucleosynthesis (BBN) can accommodate DM annihilation rate below $3 \times 10^{-26}$ cm$^3$ s$^{-1}$~\cite{GKR, KMY,GG, ADS}. Interestingly, an EMD era is a generic feature of an important class of early universe models arising from string theory constructions (for a review, see~\cite{KSW}). In these models, the modulus fields are displaced from the minimum of their potential during inflation due to misalignment~\cite{cmp}, and dominate the energy density of the post-inflationary universe due to their long lifetime. Late decay of moduli reheats the universe to temperatures below the freeze-out temperature $T_{\rm f}$, thereby rendering the WIMP miracle irrelevant in this framework. The presence of an EMD epoch in the early universe typically decreases the number of e-foldings between the horizon exit of observationally relevant cosmological perturbations and the end of inflation~\cite{LL}. However, such a change also affects inflationary predictions for the scalar spectral index $n_s$ and the tensor-to-scalar ratio $r$~\cite{EGOW,DDM}. Furthermore, \cite{EGOW} discuses some connections between supersymmetry, non-thermal DM and precision cosmology.\footnote{For some related work along this line, see \cite{related,infreh}.} This implies that cosmic microwave background (CMB) experiments may be used to constrain the non-thermal DM production from an epoch of EMD. In this paper, we explore this issue for the case with small DM annihilation rate $\langle \sigma_{\rm ann} v \rangle_{\rm f} < 3 \times 10^{-26}$ cm$^3$ s$^{-1}$. We consider various contributions to the DM abundance and obtain an absolute lower bound on the duration of the EMD era without making explicit reference to its particle physics origin. We contrast this with the upper bounds derived from the $n_s$ considerations for representative models of inflation that are compatible with the latest Planck results~\cite{planckinflation}. We show that the current~\cite{planckinflation} and future~\cite{next1} CMB experiments can tightly constrain the parameter space for an epoch of EMD. In particular, provided that the post-inflationary universe has an equation of state characterized by $w \leq 1/3$, a typical modulus-driven EMD period as the origin of non-thermal supersymmetric DM may be disfavored for inflationary models with $r \lesssim {\cal O}(0.01)$. The rest of this paper is organized as follows. In Section II, we discuss various contributions to DM production during EMD and derive an absolute lower bound on the duration of this period. In Section III, we discuss connection to inflationary observables using a parametrization of $n_s$ and $r$ that holds for a large number of inflationary models compatible with the latest observational data. In Section IV, we present our main results along with some disucusion. Finally, we conclude the paper in Section V. | We would like to emphasize that the constraints depicted in Fig.~1 are on the conservative side. First, the upper bounds on $H_{\rm dom}/H_{\rm R}$ in Eq.~(\ref{boundEMD}) are obtained assuming that the universe enters a RD phase right after inflation. Also, Eq.~(\ref{decbound}) gives a model-independent absolute lower bound on $H_{\rm dom}/H_{\rm R}$ in order to satisfy the DM relic abundance. Including model details of reheating after inflation or considering specific particle physics realizations of the EMD epoch can make the corresponding inequalities stronger and therby lead to (much) tighter constraints. We would also like to note that the CMB limits on $n_s$ used here are for the $\Lambda$CDM+$r$ model. Both the mean value and $2 \sigma$ error of $n_s$ in extensions of this model will be different, which can affect our constraints. Notably, inclusion of dark radiation results in the $2 \sigma$ allowed range $n_s = 0.9607^{+0.0176}_{-0.0168}$ for Planck data alone and $n_s = 0.9660 \pm 0.0140$ for Planck plus BK14 and BAO data~\cite{planckinflation}. The resulting $N^{\rm min}_{k_*}$ from Eq.~(\ref{Nrange}), hence the upper bounds on the duration of EMD from Eq.~(\ref{boundEMD}), will be significantly weaker in this case. Thus future CMB data~\cite{next1} will likely be needed in order to find constraints comparable to those in Fig.~1. In conclusion, we studied viability of non-thermal DM from a period of EMD in light of CMB data. Motivated by the increasingly tighter upper limits from indirect searches on the DM annihilation rate, we focused on the case with small annihilation rate $\langle \sigma_{\rm ann} v \rangle_{\rm f} < 3 \times 10^{-26}$ cm$^3$ s$^{-1}$. We found interesting constraints on the parameter space of the EMD phase by combining the lower and upper bounds on its duration from the DM relic abundance consideration and the CMB data on the scalar spectral index respectively. In particular, inflationary models with $r \lesssim {\cal O}(0.01)$ that are compatible with the latest Planck results may disfavor non-thermal supersymmetric DM from a modulus-driven EMD. This conclusion holds as long as the post-inflationary universe has an equation of state characterized by $w \leq 1/3$, and can become stronger with data from the future CMB experiments. | 18 | 8 | 1808.02659 |
1808 | 1808.00006_arXiv.txt | We present a new, computationally efficient, energy-integrated approximation for neutrino effects in hot and dense astrophysical environments such as supernova cores and compact binary mergers and their remnants. Our new method, termed ILEAS for Improved Leakage-Equilibration-Absorption Scheme, improves the lepton-number and energy losses of traditional leakage descriptions by a novel prescription of the diffusion time-scale based on a detailed energy integral of the flux-limited diffusion equation. The leakage module is supplemented by a neutrino-equilibration treatment that ensures the proper evolution of the total lepton number and medium plus neutrino energies as well as neutrino-pressure effects in the neutrino-trapping domain. Moreover, we employ a simple and straightforwardly applicable ray-tracing algorithm for including re-absorption of escaping neutrinos especially in the decoupling layer and during the transition to semi-transparent conditions. ILEAS is implemented on a three-dimensional (3D) Cartesian grid with a minimum of free and potentially case-dependent parameters and exploits the basic physics constraints that should be fulfilled in the neutrino-opaque and free-streaming limits. We discuss a suite of tests for stationary and time-dependent proto-neutron star models and post-merger black-hole-torus configurations, for which 3D ILEAS results are demonstrated to agree with energy-dependent 1D and 2D two-moment (M1) neutrino transport on the level of 10--15 percent in basic neutrino properties. This also holds for the radial profiles of the neutrino luminosities and of the electron fraction. Even neutrino absorption maps around torus-like neutrino sources are qualitatively similar without any fine-tuning, confirming that ILEAS can satisfactorily reproduce local losses and re-absorption of neutrinos as found in sophisticated transport calculations. | Neutrinos play an important role in high-energy stellar astrophysics, more precisely in the context of the birth and death of neutron stars (NSs), among others. Already in the 1990s, the neutrino-induced shock revival was theorized as a mechanism for exploding core-collapse supernovae (SNe) (\citealt{1985ApJ...295...14B}, \citealt{1990RvMP...62..801B}, see \citealt{2012ARNPS..62..407J}, \citealt{2013RvMP...85..245B} and \citealt{2015PASA...32....9F} for recent reviews). After a successful explosion, neutrinos are essential to understand the long-term cooling of the new-born NS (e.g. \citealt{2010PhRvL.104y1101H}). At the other end of their lives, NSs in binaries with another compact object (CO), either a NS, a black hole (BH) or a white dwarf (WD), may be able to merge within a Hubble time. In such scenarios, NS matter reaches very high temperatures and densities, emitting copious amounts of neutrinos (see e.g. \citealt{2011LRR....14....6S}, \citealt{2012LRR....15....8F} and \citealt{2015IJMPD..2430012R} for reviews). Despite being dynamically only a secondary ingredient, neutrinos, by their emission and absorption, drive the neutron-to-proton ratio of the NS as well as the ejected matter. This aspect will determine the distribution of synthesized elements in the ejecta \citep{2015MNRAS.452.3894G, 2015MNRAS.448..541J, 2014ApJ...789L..39W, 2015PhRvD..91f4059S, 2015ApJ...815...82L, 2016PhRvD..93d4019F, 2016PhRvD..94l3016F, 2016MNRAS.460.3255R, 2017PhRvD..96l4005B, 2018CQGra..35c4001M, 2016CQGra..33r4002L} as well as the associated electromagnetic transient, known as kilonova, powered by the radioactive decay of neutron-rich elements (\citealt{1998ApJ...507L..59L, 2005astro.ph.10256K, 2010MNRAS.406.2650M, 2011ApJ...738L..32G, 2011ApJ...736L..21R, 2016ApJ...831..190H, 2016ApJ...829..110B, 2017PhRvD..96l4005B}, see \citealt{2017ARNPS..67..253T} and \citealt{2016ARNPS..66...23F,2017LRR....20....3M} for recent reviews on r-process in NS mergers and kilonovae, respectively). Neutrinos are also needed to understand the fate of the hypermassive NS (HMNS) remnants and the evolution of the surrounding torus (if any) \citep{2014MNRAS.443.3134P, 2014MNRAS.441.3444M, 2015PhRvD..91l4021F,2017ApJ...846..114F,2017PhRvD..96l3015W}. Merger remnants have been envisioned as the central engines of short gamma-ray bursts (sGRB), and neutrino pair-annihilation has been shown to deposit significant amounts of energy, which might help in powering an ultrarelativistic jet \citep{2016ApJ...816L..30J, 2017JPhG...44h4007P}. The recent detection of a NS merger though its associated signals in the form of gravitational waves (GW) and electromagnetic radiation (EM), as predicted by theoretical models, highlights the remarkable contributions of detailed numerical simulations to our understanding of the universe. The evolution of the neutrino phase space distribution function obeys the Boltzmann transport equation. In three spatial dimensions (3D), this becomes a six-dimensional, time-dependent problem for each neutrino species, which is considerably arduous to solve when no further approximations are applied to reduce its dimensionality (e.g. \citealt{1966AnPhy..37..487L}). Therefore, numerous schemes of varying complexity and accuracy have been developed to cope with this challenging task. In the context of NS mergers, truncated moment schemes are the most sophisticated treatments successfully used. In such schemes, the angular dependence of the neutrino momentum distribution is removed by evolving a hierarchy of ("moment") equations, which are obtained from angular moment integration of the Boltzmann equation. Thus one introduces moments as angular integrals of the neutrino phase-space distribution function, i.e. neutrino energy density, flux density, pressure and higher-order moments. It is necessary, in order to close the set of equations, to find a way to express the highest employed moments, which are not evolved by but appear in the moment equations. The so-called M1 schemes (e.g. \citealt{2011PThPh.125.1255S}) time-integrate the evolution equations for the zeroth- and first-order moments, closing the system with an analytical relation expressing the highest employed moments as local functions of the evolved ones, and are used in grey (e.g. \citealt{2015PhRvD..91l4021F}) as well as energy-dependent versions (e.g. \citealt{2015MNRAS.453.3386J}). Despite their enormous current popularity for describing the transport of neutrinos, they suffer from the inability to properly handle crossing radiation beams, and thus have a tendency to overestimate the neutrino densities in the polar directions in NS-merger remnant simulations \citep{2015MNRAS.448..541J, 2018PhRvD..98f3007F}. These limitations call for alternative treatments of neutrino re-absorption, such as ray-tracing, in order to provide an accurate description of the ejecta composition, essential for predicting EM counterparts to NS mergers. Recently, an alternative Monte Carlo closure has been suggested by \cite{2018MNRAS.475.4186F}, but its applicability in merger simulations has not been demonstrated yet. Monte Carlo (MC) codes are still obviated from their direct use in full-scale merger simulations because of their tremendous computational costs and memory requirements, for which reason they have been set aside in favour of less expensive methods, except for some MC applications in snapshot calculations \citep{2015ApJ...813...38R, 2018PhRvD..98f3007F}. Ray-tracing algorithms, which solve the Boltzmann equation in one dimension, have also been employed in snapshot calculations of NS merger remnants by \cite{2015MNRAS.448..541J} in the Newtonian framework and, recently, by \cite{2018PhRvD..98j3014D} in a general-relativistic version, which also includes the effects of neutrino scattering in an approximative manner based on the existence of a precomputed M1 solution. Leakage schemes are a very popular and computationally simple approximation for the treatment of neutrino effects in NS merger simulations. They had been introduced first in Newtonian merger models as grey versions by \cite{1996A&A...311..532R, 1997A&A...319..122R, 1999A&A...344..573R, 2001A&A...380..544R} and, with a different handling of the energy integration, by \cite{2003MNRAS.342..673R, 2003MNRAS.345.1077R, 2012MNRAS.426.1940K, 2013MNRAS.430.2585R,2013RSPTA.37120272R}. More recently, leakage schemes have been used in many relativistic merger simulations, too (for example, \citealt{2011PhRvL.107e1102S,2011PhRvL.107u1101S,2012CQGra..29l4003K,2013ApJ...776...47D,2014PhRvD..90b4026F,2017CQGra..34d4002F,2014PhRvD..89j4029N,2015PhRvD..92d4045P,2016PhRvD..94d3003L,2016CQGra..33r4002L,2016PhRvD..94b4023B}). Moreover, leakage methods have been applied in evolution studies of NS merger remnants, including HMNSs \citep{2014MNRAS.443.3134P, 2015ApJ...813....2M,2018CQGra..35c4001M,2014MNRAS.441.3444M,2017MNRAS.472..904L} as well as BH-torus systems \citep{2007PThPh.118..257S,2015MNRAS.446..750F,2015MNRAS.449..390F}. The most basic version of leakage schemes accounts for neutrino energy and lepton-number losses by local source terms in the hydrodynamics equations, following the original implementations by \cite{1996A&A...311..532R} and \cite{2003MNRAS.342..673R}, possibly upgraded by gravitational redshift effects \citep{2010CQGra..27k4107S,2010CQGra..27k4103O,2013PhRvD..88f4009G}. In more advanced versions this basic functionality of leakage schemes is supplemented by treatments of a trapped neutrino component and by neutrino absorption. Such improvements have been accomplished by hybrid methods between leakage and M1 \citep{2012PTEP.2012aA304S, 2015PhRvD..91f4059S, 2016PhRvD..93l4046S,2017ApJ...846..114F,2017PhRvD..96l3012S,2018PhRvD..97b3009K} or `M0' (zeroth moment with a closure) \citep{2016MNRAS.460.3255R,2017ApJ...838L...2R,2017ApJ...842L..10R,2018ApJ...852L..29R,2018PhRvL.120k1101Z}. Alternatively, equilibrium and time-scale arguments have been used to parametrize the trapping physics, and complex propagation paths have been considered to connect the locations of neutrino production in a leakage treatment with the neutrinospheric decoupling region for describing neutrino absorption exterior to the trapping domain (see \citealt{2014MNRAS.443.3134P,2014A&A...568A..11P,2016ApJS..223...22P}). Comparisons by \cite{2015PhRvD..91l4021F, 2016PhRvD..93d4019F} reveal major differences between results with their grey M1-based SpEC code and `classical' leakage results for the first 10--15\,ms after the collision of compact binary stars. In contrast, \cite{2016ApJS..223...22P} report very good qualitative and partially quantitative agreement in key quantities when testing their Advanced Spectral Leakage (ASL) scheme against Boltzmann transport in the context of Newtonian, spherically symmetric (1D) hydrodynamic simulations of several 100\,ms of post-bounce accretion in core-collapse supernovae. However, the ASL code involves a variety of parameters that were calibrated on grounds of this considered problem. It is not obvious that the thus determined parameter values work equally well for a broader class of conditions. Moreover, also the axisymmetric (2D) and three-dimensional (3D) simulations of stellar core-collapse and post-bounce accretion, which \cite{2016ApJS..223...22P} applied their ASL code to, still contain a quasi-spherical, highly opaque neutrino source that accretes mass from the collapsing star at high rates (as in the 1D simulations). These tests are not conclusive with respect to the question how well the ASL code is able to perform in the merger case, where the remnant is rotationally deformed and not accreting. Radial profiles of the neutrino quantities for the transition from the high-opacity to the low-opacity regime and tests with non-spherical neutrino sources, which could facilitate such a judgement, are not available for upgraded leakage schemes in the literature. In this work we present a new implementation of an improved leakage treatment that does not only take into account local energy and lepton-number losses by neutrino emission, but it also accounts for the fact that neutrinos can equilibrate with matter in the optically thick regime and that they are still re-absorbed by matter when they propagate through optically thin regions. Our new method, which is termed Improved Leakage-Equilibration-Absorption Scheme (ILEAS), is designed to fulfil a number of requirements: (1) low algorithmic complexity in order to enable easy numerical realization; (2) proper and consistent reproduction of the correct physical behaviour of the neutrino-matter system at high optical depths; (3) description of the transition to the low-opacity regime with a minimum number of free parameters and ad hoc recipes of approximation; and (4) high computational efficiency that permits the calculation of large sets of merger simulations to explore the multidimensional parameter space (system masses and mass ratios, spins, orbital parameters, NS equations of state) that describes NS-NS/BH binaries. The computational efficiency is also facilitated by the fact that ILEAS, in contrast to transport calculations with explicit schemes, is not subject to any time-stepping constraints by the Courant-Friedrichs-Lewy condition (CFL). ILEAS is implemented on a 3D Cartesian grid and makes use of the grey description of the leakage loss terms applied by \cite{1996A&A...311..532R}. The greyness of the treatment benefits all of the mentioned requirements. Advancing beyond the original treatment by \cite{1996A&A...311..532R}, ILEAS introduces a new definition of the neutrino-loss time-scale based on the energy-integrated equation of flux-limited diffusion. This allows for a considerably improved description of the neutrino drain from regions of high optical depths. The effects of a trapped neutrino component are taken into account by considering neutrinos as part of an equilibrated neutrino-matter fluid in the trapping regime. Neutrino absorption in the transition to the optically thin limit is handled by a simplified ray-tracing method that adopts an analytical integration of the radiation attenuation along the ray paths of escaping neutrinos following \cite{2001A&A...368..527J}. To assess the quality of the ILEAS scheme, we consider different stages during the cooling evolution of a spherical proto-NS (PNS) and perform steady-state as well as time-dependent calculations (co-evolving the medium temperature and electron fraction on a fixed background density). We compare the leakage results for the neutrino emission with 1D neutrino transport results obtained with the ALCAR and VERTEX codes. Both of these codes are energy-dependent two-moment schemes employing an algebraic (M1) closure and a variable Eddington factor closure based on a solution of the Boltzmann equation, respectively. Moreover, we perform time-dependent calculations for the neutrino emission from optically thick (high-mass) as well as optically thin (low-mass) axisymmetric BH-accretion tori in direct comparison with ALCAR results. Our tests demonstrate very good compatibility between leakage and transport results (global quantities agree on the level of roughly 10 percent or better) with respect to radial luminosity profiles, neutrino luminosities evolving over periods of tens of milliseconds, mean energies, and spatial distributions of electron fraction and neutrino-energy absorption rates. Our paper is structured as follows. In section~\ref{sec:model} we describe the physical and algorithmic components of the ILEAS code and their numerical realization, in section~\ref{sec:tests} we present our set of neutrino transport tests for proto-NS and BH-torus models, as well as first demonstrations of the application of ILEAS in NS-NS merger models, and in section~\ref{sumary} we summarize our work. In the four following appendices~\ref{appendix:tdiffs}--\ref{appendix:nabs} we compare results for different definitions of the diffusion time-scale used in previous literature, present an overview of the neutrino opacities and source terms employed by ILEAS, discuss different versions of implementing the $\beta$-processes in the leakage treatment, and provide test results for an alternative method to compute neutrino-number re-absorption, respectively. | \label{sumary} The detection of a NS-NS merger using GW interferometers and its associated EM transients \citep{PhysRevLett.119.161101,2017ApJ...848L..12A} has opened the doors to a new era of multi-messenger astronomy. The vivid discussions about the nature of the observed kilonova (e.g. \citealt{2017ApJ...848L..17C, 2017ApJ...848L..18N, 2017ApJ...848L..19C, 2017Natur.551...80K, 2017Natur.551...75S, 2017Natur.551...67P, 2017ApJ...850L..37P, 2018MNRAS.481.3423W}) has brought to light the need of a reliable understanding not only of the composition of the ejected material, but also its dependence on the direction of ejection. At the same time, the error-bars associated with the measured NS masses, together with the underlying uncertainty of the EoS of NS matter, call for the exploration of a wide distribution of initial conditions for numerical simulations. In order to reconcile these two needs we have introduced \textsc{ILEAS}, an improved leakage scheme which accounts for the basic physical effects of neutrino transport at a moderate computational cost. \textsc{ILEAS} is ideal for exploring wide parameter spaces in three dimensions, where $\sim$10 per cent of accuracy is enough to capture the essential impact of weak interactions. Leakage models have been used to emulate neutrino losses in the context of NS mergers since the 1990's \citep{1996A&A...311..532R}, but little work has been devoted to assess their accuracy (exceptions are \citealt{2016PhRvD..93d4019F,2016ApJS..223...22P}). In fact, we show that, in their standard formulations, leakage schemes have a tendency to overproduce neutrinos in the region close to the neutrinosphere, which could lead to numerical artefacts in near-surface regions and incorrect estimates of the ejecta composition. Moreover, the traditional leakage schemes were only a simple ansatz to estimate the local neutrino losses, ignoring other important physical effects inherent to neutrino transport, such as equilibration or re-absorption. In the recent years, truncated moment schemes have been developed and successfully used in the context of NS mergers, providing a more sophisticated alternative to leakage schemes. These approaches, however, also possess disadvantages and shortcomings of their own, such as problems with crossing flows (see, e.g., \citealt{2018PhRvD..98f3007F}) or the need of more computational resources, in particular when combined with a ray-tracing scheme for computing high-resolution neutrino distributions. While not a proper neutrino transport scheme, \textsc{ILEAS} is able to capture all the aforementioned physical effects, yet retaining the simple and inexpensive aspects of leakage schemes. Its improved diffusion time-scale, obtained directly from the flux-limited diffusion equation, provides a much better estimate of the neutrino losses in optically thick regions. This is reinforced by the inclusion of equilibration: the equilibration step ensures the recovery of the correct lepton fractions in the $\beta$-equilibrium regime, and the EoS also includes the energy and pressure contributions of the trapped neutrinos. Finally, by means of a simple multi-dimensional ray-tracing algorithm, we account for the re-absorption in optically thin conditions of neutrinos leaking out from the system. In order to keep the absorption module computationally efficient, we decided to resort to a grey approximation for \textsc{ILEAS}. However, our spectral calculation of the diffusive flux allows us to approximately capture the energy-dependent decoupling of neutrinos from matter in the integrated diffusion time-scale. Our results show that, despite the inherent approximations, \textsc{ILEAS} is sufficiently good to reproduce the results of more sophisticated transport schemes on the level of 10 per cent, locally and globally. Motivated by its future application in the context of NS mergers, we tested the performance of \textsc{ILEAS} by comparison to available simulations representing some of the typical conditions encountered during NS mergers. We presented the results obtained with \textsc{ILEAS} applied on 3D mappings of several PNS cooling snapshots from the 1D \textsc{VERTEX-PROMETHEUS} simulation (Sr) with energy-dependent neutrino transport performed by \cite{2010PhRvL.104y1101H}. For all tested snapshots, ranging from 0.2~s until 1.5~s post-bounce, \textsc{ILEAS} was able, after a short relaxation of the medium (i.e., after evolving temperature and $Y_e$ for 5\,ms), to reproduce not only the total \textsc{VERTEX} luminosities, but also the complete radial luminosity profiles within $\sim$10 per cent accuracy. In order to provide a more detailed comparison, we also tested \textsc{ILEAS} on a snapshot obtained from the evolution of the same PNS performed by \textsc{ALCAR} \citep{2015MNRAS.453.3386J}, which includes an energy-dependent M1 transport solver and exactly the same neutrino reactions as \textsc{ILEAS} for $\nu_e$ and $\bar{\nu}_e$. As with the \textsc{VERTEX} cases, \textsc{ILEAS} reached an agreement within 10 per cent accuracy with the \textsc{ALCAR} results. Furthermore, we evolved the temperature and $Y_e$ of the \textsc{ALCAR} snapshot (keeping the density fixed) for 50~ms with both \textsc{ALCAR} and \textsc{ILEAS}, and the good agreement was maintained throughout the simulation. As possible remnants of CO mergers, BH-torus systems provide a useful scenario for testing the performance of our scheme in the low optical depth limit. Snapshot calculations allowed us to attest the capability of our absorption treatment to capture the qualitative features of neutrino absorption in comparison with the results attained by \textsc{ALCAR}. Furthermore, we evolved two BH-torus models for 50~ms using \textsc{ALCAR} and \textsc{ILEAS}, in the same fashion as the PNS snapshot. In spite of the initial over-cooling caused by the unavoidable transient produced when switching on \textsc{ILEAS}, the neutrino luminosities of both models preserved an agreement of $\sim$10 per cent. In conclusion, \textsc{ILEAS} has been shown to reproduce within $\sim$10 per cent accuracy basic results of more sophisticated transport schemes also in multi-dimensional scenarios. Albeit not as accurate as full-fledged 3D transport, \textsc{ILEAS} includes all the relevant physical effects of neutrino transport and surpasses the quality of previous, conventional leakage schemes while retaining most of their efficiency and simplicity. These features make \textsc{ILEAS} an appropriate description of neutrinos for numerical simulations of NS mergers, where the relatively short evolution time-scales may not require a full-scale 3D neutrino transport to still obtain a consistent picture of the composition evolution of merger medium and ejecta. The exploration of the vast parameter space of possible binary configurations demands computationally efficient but sufficiently accurate codes. \textsc{ILEAS} is intended to serve this requirement. In section~\ref{sec:tests:NSNS} we presented first NS-NS merger simulations for two different EoSs employing \textsc{ILEAS} for the neutrino treatment. The models demonstrate the feasibility, numerical stability and basic agreement with results available in the literature, although detailed comparisons of the neutrino quantities are not possible because of the combination of different neutrino treatments and different hydrodynamics schemes. Nevertheless, we caution the reader that \textsc{ILEAS} cannot be perceived as a perfect replacement for a full neutrino transport scheme, especially not in situations where the transfer of energy and leptons from one location to another by neutrino diffusion is crucial to describe the interior evolution of an object, as, for example, in the case of the long-time neutrino-cooling of PNSs. \textsc{ILEAS} is a suitable alternative to neutrino transport codes particularly in conditions where the dynamical changes of the system happen on a time-scale shorter than or comparable to the neutrino diffusion time-scale. Examples are the rapid evolution of two merging NSs or the cooling of (semi-)transparent tori around merger remnants. Finally, for its use in consistent GR simulations, \textsc{ILEAS} needs to be extended to include effects such as Doppler shift, time retardation and general relativistic gravitational ray bending, which were not essential in the test calculations presented in this work. In future versions, \textsc{ILEAS} will also be supplemented by a module describing the effects of neutrino-antineutrino annihilation. A computationally considerably more expensive alternative to ILEAS is advertised by \cite{2018PhRvD..98f3007F}, who plan to upgrade their grey M1 scheme by an Eddington tensor closure obtained from a Monte Carlo (MC) solution of the Boltzmann equation exterior to regions of high optical depths, instead of using the analytical closure relations applied so far. Whether such a hybrid code yields stable results with acceptable numerical ``noise'' for affordable low-resolution MC calculations will have to be demonstrated. For reasons of accuracy, \cite{2018PhRvD..98f3007F} recommend to evolve the neutrino number densities as well as the neutrino energy densities in order to obtain reasonably accurate local estimates of the average neutrino energies. In order to evaluate errors of M1 results with analytical closure relations, they performed a time-dependent calculation over 4.5\,ms with their relativistic MC solver for a HMNS as a representative remnant of a NS-NS merger, using the time-dependent fluid quantities from the M1 radiation-hydrodynamics run and not feeding back the MC results into fluid or M1 transport solutions. They found relative differences of 10--30 per cent in the average neutrino energies between the M1 and the MC transport results and concluded that this implies that the absorption and scattering opacities can be off by $\sim$30 per cent up to close to a factor of 2, dependent on positions closer to the polar axis (i.e.\ the rotation axis of the remnant) or farther away from it. Moreover, because of artificial shocks associated with the use of a non-linear, algebraic closure relation, they diagnosed that the M1 code accumulates neutrinos close to the polar axis, leading to an excess of the neutrino density in the polar regions by about 50 per cent for $\nu_e$ and $\bar\nu_e$ and by nearly a factor of 2 for $\nu_x$. The $\nu\bar\nu$ pair-annihilation rate above the poles of the HMNS is underestimated by factors of 2--3 by the M1 description. Similar results had already been reported in the context of BH-torus systems by \cite{2015MNRAS.448..541J} (in the appendix there), who also found an overestimation of the neutrino densities in the polar regions when comparing their energy-dependent M1 scheme, ALCAR, with a ray-tracing Boltzmann solution. However, in the tested BH-torus scenario, the results by \cite{2015MNRAS.448..541J} imply an \textit{over}-estimation of the pair-annihilation rate around the polar axis. While the spectral and opacity differences may be handled better by fully energy dependent (and considerably more complex and costly) transport codes such as ALCAR, the overestimated number densities and underestimated pair-annihilation rates have to await their cure through a replacement of the analytic closure by a Boltzmann transport solution, possibly based on MC results. In view of the considerable error margins associated with a grey M1 approximation and considering the high computational demands of future hybrid schemes, our ILEAS method constitutes itself as an interesting option for the next generation of NS-NS/BH merger simulations surveying the huge multi-dimensional parameter space of possibilities. ILEAS is not only computationally very efficient but also appears to be competitive concerning its accuracy compared to other forefront developments of neutrino transport treatments for CO mergers and their remnants. The work by \cite{2018PhRvD..98f3007F} underlines that M1 solutions have their own shortcomings when applied to the highly aspherical environments of merger remnants. For this reason our tests with BH-torus systems, comparing ILEAS to M1 results from the ALCAR code, cannot be considered as finally conclusive regarding the accuracy of ILEAS. Direct comparisons of ILEAS and MC results would be desirable. | 18 | 8 | 1808.00006 |
1808 | 1808.00230_arXiv.txt | Massive stars are rarely seen to form in isolation. It has been proposed that association with companions or clusters in the formative stages is vital to their mass accumulation. In this paper we study IRAS~18144-1723, a massive young stellar object (YSO) which had been perceived in early studies as a single source. In the CO(3-2) line, we detect an outflow aligned well with the outflow seen in H$_2$ in this region. We show that there are at least two YSOs here, and that the outflow is most likely to be from a deeply embedded source detected in our infrared imaging. Using multi-wavelength observations, we study the outflow and the embedded source and derive their properties. We conclude that IRAS~18144 hosts an isolated cloud, in which at least two massive YSOs are being born. From our sub-mm observations, we derive the mass of the cloud and the core hosting the YSOs. | Recent studies suggest that the primary mechanism for the formation of massive stars is disk accretion as in their low-mass counterparts \citep{Arce07, Beuther02a, Varricatt10, Davis04}. Collimated outflows discovered from massive star forming regions in these studies indicate the presence of accretion discs. The cavities carved by the protostellars outflow provide a path for the radiation to escape, thereby reducing the radiation pressure on the accreted matter and allowing the mass accumulation in massive YSOs through accretion \citep{Krumholz05}. However, unlike low-mass stars, massive stars are rarely observed to form in isolation. Instead, they are seen to be associated with companions or clusters suggesting that such associations may play a prominent role in their formation \citep{Lada03}. Massive YSOs are located in the galactic plane at large distances from us, so we need observations at infrared and longer wavelengths at high angular resolution to understand their formation. With the availability of large telescopes operating at these wavelengths, many of the massive YSOs, which appeared to form in isolation in older studies, are now being resolved into binaries or multiples (e.g. \citet{Varricatt10}). In this paper, we present a detailed observational study of the massive YSO IRAS~18144-1723 (hereafter IRAS~18144). IRAS~18144 is located near the galactic plane (l=13.657$^{\circ}$, b=-0.6$^{\circ}$), and is associated with a dense core of far-IR luminosity 1.32$\times$10$^4$~L$_\odot$, detected in NH$_3$ emission \citep{Molinari96}. Using the line-of-sight radial velocity of the NH$_3$ emission (+47.3\,km~s$^{-1}$), they derived a kinematic distance of 4.33\,kpc. Water and methanol masers are also detected from this source \citep{Palla91, Szymczak00, Kurtz04, Gomez-Ruiz16}. \citet{Molinari98} did not detect 6-cm radio emission associated with the IRAS source sugessting that it is in a pre-UCH{\sc{ii}} stage. Observations by \citet{Zhang05} did not detect any CO(2--1) outflow towards this region. However, through near-IR imaging in the H$_2$ line at 2.1218\,$\mu$m and in the $K$ filter, \citet{Varricatt10} discovered an E--W jet appearing to emanate from a bright near-infrared source (referred to as `A') located at $\alpha$=18:17:24.38, $\delta$=-17:22:14.7\footnote{All coordinates given in this paper are in J2000}. We further observed this region at multiple wavelengths from near-IR to sub-mm in order to uncover the nature and possible multiplicity of source `A'. The new observations are analysed along with archival data. \section[]{Observations and data reduction} \subsection[]{Near-IR imaging with WFCAM} We observed IRAS~18144 using the United Kingdom Infrared Telescope (UKIRT), Hawaii, and the UKIRT Wide Field Camera (WFCAM; \cite{Casali07}). WFCAM employs four 2048$\times$2048 HgCdTe HawaiiII arrays, each with a field of view of 13.5\arcmin$\times$13.5\arcmin at an image scale of 0.4\arcsec~pixel$^{-1}$. Observations were obtained in the the near-IR $J, H$ and $K$-band filters, and in narrowband filters centred on the H$_2$ $\upsilon$ = 1--0 S(1) line at 2.1218~$\mu$m and the [Fe {\sc{II}}] $a^4 D_{7/2} - a^4 F_{9/2}$ forbidden emission line at 1.6439~$\mu$m. H$_2$ and [\ion{Fe}{ii}] are good tracers of star formation activity. H$_2$ emission may be fluorescently excited in photodissociation regions associated with massive and intermediate mass young stars or in planetary nebulae, or shock excited in jets and warm entrained molecular outflows from YSOs. [Fe {\sc{II}}] is often used as a tracer of ionized region at the base of the jet \citep{Davis11}. \begin{table} \centering \caption{Imaging observations using WFCAM and UIST} \label{tab:UKIRT_imaging_log} \begin{tabular}{@{}lllll@{}} \hline UTDate & Filters & Exp. time &Int. time &FWHM \\ yyyymmdd & & (sec) &(sec) &(arcsec) \\ \hline \multicolumn{5}{c}{WFCAM}\\ \hline 20140620 & $J$ & 10,2 &360, 40 & 1.11, 1.10 \\ 20140620 & $H$ & 5,1 &180, 20 & 1.10, 1.21 \\ 20140620 & $K$ & 5,1 &180, 20 & 1.05, 0.91 \\ 20140620 & 1-0S1 & 40 &1440 & 1.07 \\ 20140621 & 1-0S1 & 40 &2$\times$1440 & 0.82, 0.79 \\ 20140621 & $J$ & 10 &360 & 0.84 \\ 20140621 & $H,K$ & 5,5 &180, 180 & 0.78, 0.77 \\ 20140622 & 1.644FeII & 40 &1440 & 0.79 \\ 20140623 & 1.644FeII & 40 &1440 & 0.86 \\ 20170522 & $J$ & 5 &720 & 0.99 \\ 20170522 & $H, K$ & 5, 5 &360, 180 & 0.82, 0.84 \\ 20180330 & $J$ & 5 &720 & 0.94 \\ 20180330 & $H, K$ & 5, 5 &360, 180 & 0.88, 0.87 \\ \hline \multicolumn{5}{c}{UIST}\\ \hline 20140820 & $L'$ &0.4$\times$50$^{\mathrm{a}}$ &160 &0.48 \\ 20140820 & $M'$ &0.175$\times$50 &105 &0.5 \\ \hline \end{tabular} \begin{list}{}{} \item {$^{\mathrm{a}}$exp. time $\times$ coadds} \end{list} \end{table} \begin{figure*} \centering \includegraphics[width=17.7cm,clip]{IRAS18144_JHH2_cs.pdf} \caption{Left: A $2\arcmin\times2\arcmin$ field of the $J$ (blue), $H$ (green), $H_2$ (red) colour composite image of IRAS~18144 observed using WFCAM. Right: The continuum-subtracted H$_2$ image of the same field. The blue (continuous) and red (dashed) contours are generated from the CO(3--2) intensity maps of the blue- and red-shifted lobes of the outflow detected in our JCMT+HARP observations, integrated in velocity ranges of 23.3 -- 44.0\,km~s$^{-1}$ and 50.6 -- 71.3\,km~s$^{-1}$ respectively. The H$_2$ emission knot MHO~2302 and the outflow source candidate (`A') identified by \citet{Varricatt10} are labelled.} \label{Fig:JHH2_cs} \end{figure*} The observations were performed by dithering the object to nine positions separated by a few arcseconds and using a 2$\times$2 microstep, resulting in a pixel scale of 0.2\arcsec~pixel$^{-1}$. The sky conditions were clear. Table \ref{tab:UKIRT_imaging_log} gives a log of the WFCAM observations. Preliminary reduction of the data was performed by the Cambridge Astronomical Survey Unit (CASU). The photometric system and calibration are described in \cite{Hewett06} and \cite{Hodgkin09} respectively. The pipeline processing and science archive are described in \cite{Irwin04} and \cite{Hambly08}. Further reduction was carried out using the Starlink packages {\sc{kappa}} and {\sc{ccdpack}}. The left panel of Fig. \ref{Fig:JHH2_cs} shows a {\em J,H,H$_2$} ($J$-blue, $H$-green, $H_2$-red) colour composite image in a 2$\arcmin\times$2$\arcmin$ field centred on IRAS~18144. The long-integration $J$ and $H$ images obtained on UT 20140620 and 20140621 and both $H_2$ images obtained on 20140620 were averaged for constructing Fig. \ref{Fig:JHH2_cs}. The narrow-band H$_2$ and [Fe {\sc{ii}}] images were continuum subtracted using scaled {\em K} and {\em H} images respectively observed closest in time as well as with good agreement in seeing, adopting the procedure given in \citet{Varricatt10}. The right panel of Fig. \ref{Fig:JHH2_cs} shows the continuum-subtracted H$_2$ image. The extended emission features are continuum subtracted well. The point sources show positive or negative residuals due to the difference in seeing between the two images or due to reddening. The H$_2$ emission knot (MHO~2302) and the outflow source candidate (`A') identified by \citet{Varricatt10} are labelled. No line emission was detected in the [Fe {\sc{ii}}] image, so it is not shown here. Source `A' of \citet{Varricatt10} is detected well in our $K$-band image as a highly reddened object embedded in nebulosity (see \S{\ref{cs_nc}}). Its detection is marginal in $H$, and most of the emission is from the nebulosity. It is not detected in $J$; only the nebulosity is seen. \subsection[]{$L'$ and $M'$ imaging using UIST} $L'$ (3.77\,$\mu$m) and $M'$ (4.69\,$\mu$m) imaging observations were obtained using the UKIRT 1--5\,$\mu$m Imager Spectrometer (UIST, \citet{Ramsay04}). UIST employs a 1024$\times$1024 InSb array. The 0.12\arcsec~pixel$^{-1}$ camera of UIST was used, which has an imaging field of 2\arcmin$\times$2\arcmin\ per frame. The observations were performed by dithering the object to four positions on the array, separated by 20\arcsec\ each in RA and Dec from the base position. Each pair of observations was flat-fielded and mutually subtracted, resulting in a positive and negative beam. The final mosaic was constructed by combining the positive and negative beams. Table \ref{tab:UKIRT_imaging_log} gives the details of the observations. The data reduction was carried out using the facility pipeline {\textsc{oracdr}} \citep{cavanagh08} and using {\sc{kappa}} and {\sc{ccdpack}}. Astrometric calibration was done using the 2MASS positions of the objects detected in our images. The UKIRT standard star GL748 was observed with the same dither pattern for photometric calibration. Fig. \ref{Fig:18144_UIST} shows a 25\arcsec$\times25\arcsec$ section of the $M'$ image. Source `A' was detected well in $L'$ band as a single object. In $M'$, we detect a deeply embedded source $\sim$2.6\arcsec\ NE of `A'; it is labelled `B' on the figure. The magnitudes for sources `A' and `B' derived from our images are given in Table \ref{tab:UIST_Mich_flx}. \begin{table} \caption{Results from UIST imaging} \label{tab:UIST_Mich_flx} % \centering % \begin{tabular}{lllll} % \hline\hline \\[-2mm] Src &RA &Dec &\multicolumn{2}{c}{magnitudes$^{\mathrm{b}}$} \\ ID & & &$L'$ &$M'$ \\ \hline \\[-2mm] A &18:17:24.375 &-17:22:14.71 &6.54 (0.04) &5.57 (0.07) \\ B &18:17:24.239$^{\mathrm{a}}$ &-17:22:12.87$^{\mathrm{a}}$ &Not det. &9.57 (0.14) \\ \hline \end{tabular} \begin{list}{}{} \item $^{\mathrm{a}}$Derived from the Michelle image after matching the coordinates of source `A' derived from the $K$-band image; $^{\mathrm{b}}$A 4\arcsec-diameter photometric aperture was used. `B' being faint, a 1.2\arcsec-diameter aperture was used and aperture correction was applied. The values given in parenthesis are the 1-$\sigma$ errors in photometry. \end{list} \end{table} \begin{figure} \centering \includegraphics[width=8.4cm,clip]{IRAS18144_M_25as.pdf} \caption{A $25\arcsec\times25\arcsec$ field of the $M'$ image of IRAS~18144 observed using UKIRT and UIST. The blue (continuous) and red (dashed) contours are generated from the integrated CO(3--2) maps of the red- and blue-shifted lobes of the outflow. } \label{Fig:18144_UIST} \end{figure} \subsection[]{Michelle imaging} Michelle \citep{Glasse97} is a mid-IR imager/spectrometer at UKIRT, employing an SBRC Si:As 320$\times$240-pixel array. It has a field of view of 67.2\arcsec$\times$50.4\arcsec with an image scale of 0.21\arcsec~pix$^{-1}$. We observed IRAS~18144 using Michelle on multiple nights in four filters centered at 7.9, 11.6, 12.5 and 18.5\,$\mu$m. The 7.9\,$\mu$m filter has a 10\% passband, and the other filters have 9\% passbands. The sky conditions were photometric with the atmospheric opacity at 225\,GHz measured with the CSO (Caltech Submillimeter Observatory) dipper ($\tau_{225GHz}$) was $\sim$0.07. BS~6705 and BS~7525 were used as the standard stars on 20040329~UT, and BS~6869 was the standard for the rest of the observations. Table \ref{tab:Mich_imaging_log} shows the details of the observations. The details of the data reduction are the same as in \citet{Varricatt13}. Astrometric corrections were applied by adopting the position of source `A' from \citet{Varricatt10}. Fig. \ref{Fig:18144_Mich} shows a colour composite image in a 25\arcsec$\times$25\arcsec\ region constructed from our Michelle images at 18.5, 12.5 and 7.9\,$\mu$m. Aperture photometry was performed using an aperture of diameter 14 pixels (2.94$\arcsec$). The flux densities measured for sources `A' and `B' are given in Table \ref{tab:Mich_imaging_log}. \begin{figure} \centering \includegraphics[width=8.4cm,clip]{IRAS18144_Mich_25as_mod.pdf} \caption{ A colour composite image (18.5\,$\mu$m (red), 12.5\,$\mu$m (green), 7.9\,$\mu$m (blue)) of a $25\arcsec\times25\arcsec$ field from our Michelle images. The contours are generated from the $K$-band image presented in \citet{Varricatt10}, and are at arbitrary levels. } \label{Fig:18144_Mich} \end{figure} \begin{table*} \centering \begin{minipage}{140mm} \caption{Observations using UKIRT and Michelle} \label{tab:Mich_imaging_log} \begin{tabular}{@{}llllllll@{}} \hline UTDate & Filter & Exp. time &Coadds &Int. time &FWHM &\multicolumn{2}{c}{Flux density in Jy$^{\mathrm{a}}$}\\ & & (sec) & &(sec) &(arcsec) &Source `A'&Source `B'\\ \hline 20151107 & 7.9\,$\mu$m &0.08 &155 &148.8 &0.73 &2.17 (0.16) &0.85 (0.1) \\ 20160826 & 7.9\,$\mu$m &0.08 &155 &248.0 &0.72 &2.02 (0.04) &0.92 (0.1) \\[3mm] 20151107 & 11.6\,$\mu$m &0.09 &144 &155.52 &0.75 &2.23 (0.04) &0.40 (0.05)\\ 20160826 & 11.6\,$\mu$m &0.09 &144 &259.2 &0.76 &1.93 (0.04) &0.40 (0.12)\\[3mm] 20040329 & 12.5\,$\mu$m &0.08 &155 &99.2 &0.82 &3.16 (0.18) &0.96 (0.13)\\ 20160826 & 12.5\,$\mu$m &0.07 &168 &141.12 &0.83 &2.55 (0.06) &0.93 (0.06)\\ 20160920 & 12.5\,$\mu$m &0.07 &168 &161.28 &0.87 &2.74 (0.13) &0.93 (0.07)\\[3mm] 20151107 & 18.5\,$\mu$m &0.03 &338 &162.24 &1.08 &5.27 (0.09) &2.41 (0.13)\\ 20160826 & 18.5\,$\mu$m &0.03 &338 &283.92 &1.08 &4.54 (0.04) &2.25 (0.06)\\ 20160920 & 18.5\,$\mu$m &0.03 &338 &162.24 &1.02 &5.63 (0.36) &2.28 (0.24)\\ \hline \end{tabular} \begin{list}{}{} \item $^{\mathrm{a}}$The values given in parenthesis are the 1-$\sigma$ errors in the flux density measured in the four beams from the Michelle image observed with chopping and nodding. \end{list} \end{minipage} \end{table*} \subsection{Archival imaging data} The field containing IRAS~18144 was covered in the {\it Spitzer} GLIMPSE survey using the Infrared Array Camera \citep[IRAC;][]{Fazio04}, and the MIPSGAL survey using MIPS \citep{Rieke04}. We downloaded the IRAC images and catalogue and the MIPS images from the {\it Spitzer} Heritage Archive\footnote{http://sha.ipac.caltech.edu/applications/Spitzer/SHA/}. Source `A' was detected well in the IRAC bands 1--4 (centered at 3.5, 4.5, 5.8 and 8.0\,$\mu$m respectively). `B' was not detected in bands 1 and 2, and was only marginally detected in band 3. It was detected well in band 4, but the psf merges with that of `A'. The {\it Spitzer} catalog lists its location as ($\alpha$=18:17:24.25, $\delta$=-17:22:13.16), which agrees with our position (Table \ref{tab:UIST_Mich_flx}) within the 0.3\arcsec\ error in the coordinates given in the GLIMPSE catalogue. IRAC 5.8 and 8.0\,$\mu$m flux densities listed for this source are 264.0$\pm$30.2 and 534.4$\pm$56.5\,mJy respectively. However, note that in UIST $M'$ (4.7\,$\mu$m), where `B' is resolved well from `A', source `B' has a flux density of only 24.2$\pm$5\,mJy (9.57 mag.; Table \ref{tab:UIST_Mich_flx}). The flux estimates of `B' are likely to be affected by the proximity of `A'. Therefore, we do not use the 5.8\,$\mu$m flux of `B' in our analysis. The source is detected in the MIPS band 1 (23.68\,$\mu$m) images at ($\alpha$=18:17:24.24, $\delta$=-17:22:12.61), but it is saturated making the photometry unusable. At a 2.45\arcsec~pix$^{-1}$ spatial resolution of the MIPS image, `A' and `B' are not resolved, but it should be noted that the centroid is much closer to `B' than to `A'. IRAS~18144 was observed in the WISE \citep{Wright10} and AKARI \citep{Murakami07} surveys. The sources `A' and `B' are not resolved by WISE. Our imaging at higher angular resolution shows that the emission from `A' will the dominant source in the WISE bands W1, W2 and W3 at 3.4, 4.6, 12\,$\mu$m respecitvely, so the magnitudes measured in those bands are not used in our analysis. Table \ref{tab:Mich_imaging_log} shows that `A' contributes $\sim$70\% of the total flux at 18.5\,$\mu$m, so we use the band W4 (22\,$\mu$m) magnitude of -1.191$\pm$0.011 as only an upper limit in the SED analysis. The flux densities measured in the AKARI-IRC 9 and 18\,$\mu$m bands were not used in this study as they will be dominated by source `A', and our Michelle observations at better spatial resolution cover these wavelengths. AKARI-FIS 90 and 140\,$\mu$m flux densities were also not used as they will be affected by the large bandwidth of these filters. The FIS 65 and 160$\mu$m flux densities were used after applying the colour correction given in \citet{Yamamura10}. We used the Herschel PACS \citep{Poglitsch10} (70 and 160\,$\mu$m), and SPIRE \citep{Griffin10} (250\,$\mu$m, 350\,$\mu$m and 500\,$\mu$m) level\,2.5 maps (observation ID \#1342218999 and \#1342219000) publicly available from the Herschel Science Archive. `A' and `B' are not resolved in the Herschel maps. Aperture photometry of the source was performed using HIPE and aperture corrections were applied. Table \ref{tab:Herschel_flx} gives the flux densities for sources `A' and `B' combined, measured from the PACS and SPIRE maps. \begin{table} \centering \begin{minipage}{140mm} \caption{Flux density of `A+B' from the Herschel maps} \label{tab:Herschel_flx} \begin{tabular}{@{}llll@{}} \hline Wavelength &Aperture &Flux density &Error\\ ($\mu$m) &diameter (arcsec) & (Jy) & (Jy) \\ \hline 70 &24 &525.6 &20.5 \\ 160 &44 &938.6 &27.7 \\ 250 &44 &454.5 &18.8 \\ 350 &60 &191.7 &12.4 \\ 500 &84 &74.7 &2.2 \\ \hline \end{tabular} \end{minipage} \end{table} \subsection{Near-IR spectroscopy using UIST} We obtained near-IR spectroscopic observations of IRAS~18144 using UKIRT and UIST on UT\,20140824. The sky conditions were clear with good seeing ($\sim$0.24\arcsec\ in the $K$ band). The HK grism was used along with a 4-pix-wide (0.48\arcsec) slit, giving a wavelength coverage of 1.395--2.506\,$\mu$m, and spectral resolution of $\sim$500. Flat field observations were obtained prior to the target observations by exposing the array to a black body mounted on the instrument. For wavelength calibration, the array was exposed to an Argon arc lamp. An early-type telluric standard was also observed at the same airmass as that of the target. We adopted a slit angle of 85.2$^{\circ}$ West of North so that source `A' and the H$_2$ emission feature MHO\,2302 are both on the slit. For the target field, we nodded the telescope between the source position and a blank field nearby to enable good sky subtraction. The total on-chip exposure time for the target was 600\,sec. The data reduction was carried out using the UKIRT pipeline {\sc{oracdr}} and Starlink {\sc{kappa}} and {\sc{figaro}} \citep{currie08}. The spectra from the two nodded beams of the reduced spectral image of standard star were extracted, averaged and wavelength calibrated. It was then divided by the spectrum of a blackbody of temperature similar to its photosperic temperature, and the photospheric absorption lines were interpolated out. The flat-fielded and sky-subtracted target spectal image was wavelength calibrated and divided by the interpolated standard star spectrum. The ratioed spectrum was then flux calibrated using the magnitudes of source `A' derived from our photometry. Fig. \ref{Fig:UIST_spim} shows the calibrated spectral image of IRAS~18144. The extracted spectrum of source `A' is shown in the upper panel of Fig. \ref{Fig:UIST_sp_csjet}. The most prominent feature in the spectrum of source `A' is the Br$\gamma$ emission line. The emission from MHO\,2302 is composed purely of the ro-vibration lines of H$_2$. The spectrum of MHO\,2302 extracted in 16 rows (7.68\,arcsec$^2$; the region shows in the dashed box in Fig. \ref{Fig:UIST_spim}) is shown in the lower panel of Fig. \ref{Fig:UIST_sp_csjet}. The integrated intensities of the H$_2$ lines measured from this spectrum are given in Table \ref{tab:H2_lines}. \begin{figure*} \centering \includegraphics[width=17.7cm,clip]{IRAS18144_spim_20140824.pdf} \caption{UIST spectral image of IRAS~18144 containing source `A' and the outflow lobe MHO\,2302. The image covers a wavelength regime of 14480--25059~\AA\ on the X axis and has a spatial extent of 221 pixels (26.5$\arcsec$) on the Y axis. The dashed box show the region within which the spectrum of MHO\,2302 is extracted } \label{Fig:UIST_spim} \end{figure*} \begin{figure*} \centering \includegraphics[width=17.7cm,clip]{IRAS18144_sp_csjet.pdf} \caption{The upper panel shows the spectrum of source `A'. The lower panel shows the spectrum of MHO\,2302, integrated over a field of 7.68\,arcsec$^2$ (the region enclosed by the dashed box drawn in Fig. \ref{Fig:UIST_spim}). } \label{Fig:UIST_sp_csjet} \end{figure*} \begin{table} \caption{H$_2$ line intensities} \label{tab:H2_lines} % \centering % \begin{tabular}{lllllll} % \hline\hline \\[-2mm] $\lambda_{air}$& Species&Intensity &Error &Einstein &g$^{\mathrm{a}}$ &E($\upsilon$,J) \\ (\AA) & & W/m$^2$ &W/m$^2$ &A (s$^{-1}$) & &(K) \\ \hline 22235& 1--0 S0& 3.51E-18 & 4.00E-20 &2.53E-7 & 5 & 6470 \\ 21218& 1--0 S1& 1.34E-17 & 2.00E-19 &3.47E-7 & 21 & 6960 \\ 20338& 1--0 S2& 5.10E-18 & 5.00E-20 &3.98E-7 & 9 & 7580 \\ 19576& 1--0 S3& 1.05E-17 & 2.00E-19 &4.21E-7 & 33 & 8370 \\ 17480& 1--0 S7& 1.65E-18 & 1.50E-19 &2.98E-7 & 57 & 12800 \\ 22477& 2--1 S1& 2.06E-18 & 4.00E-20 &4.98E-7 & 21 & 12600 \\ 24066& 1--0 Q1& 1.27E-17 & 2.00E-19 &4.29E-7 & 9 & 6150 \\ 24134& 1--0 Q2& 5.40E-18 & 2.00E-19 &3.03E-7 & 5 & 6470 \\ 24237& 1--0 Q3& 1.24E-17 & 3.00E-19 &2.78E-7 & 21 & 6960 \\ 24375& 1--0 Q4& 4.78E-18 & 1.00E-19 &2.65E-7 & 9 & 7580 \\ 24548& 1--0 Q5& 6.43E-18 & 2.00E-19 &2.55E-7 & 33 & 8370 \\ 24756& 1--0 Q6& 4.21E-18 & 1.50E-19 &2.45E-7 & 13 & 9290 \\ 25001& 1--0 Q7& 8.88E-18 & 5.00E-19 &2.34E-7 & 45 & 10300 \\ \hline \end{tabular} \begin{list}{}{} \item $^{\mathrm{a}}$ Statistical weight \end{list} \end{table} \subsection{Sub-mm continuum observation using SCUBA-2} We observed IRAS~18144 using the James Clerk Maxwell Telescope (JCMT) and Submillimetre Common-User Bolometer Array 2 (SCUBA-2; \cite{Holland13}) on UT~20140626. SCUBA-2 maps simulteneously at 450 and 850\,$\mu$m. Each band uses four 32$\times$40 Transition Edge Sensor (TES) arrays, with a main beam size of 7.9\arcsec\ and 13.0\arcsec\ respectively at 450\,$\mu$m and 850\,$\mu$m. The observation was performed by moving the telescope in a ``CV Daisy'' pseudo-circular pattern\footnote{See http://www.eaobservatory.org/jcmt/instrumentation/\\continuum/scuba-2/ for more details about observatios using SCUBA-2}. A field of diameter of $\sim$12\arcmin\ was obtained. The weather conditions were good with $\tau_{225\,GHz}$=0.08 during the observations. Pointing and extinction corrections were applied using pointing and flux standards observed before the target observations. The data reduction was performed using the Starlink package {\sc{smurf}}\footnote{Also see http://www.starlink.ac.uk/docs/sc21.htx/sc21.html}. The maps were sampled down to 2\arcsec~pix$^{-1}$ at 450\,$\mu$m and 4\arcsec~pix$^{-1}$ at 850\,$\mu$m. We reach RMS noise levels of 2.04\,mJy\,arcsec$^{-2}$ (213\,mJy/beam) and 0.031\,mJy\,arcsec$^{-2}$ (7\,mJy/beam) respectively at the centers of the 450\,$\mu$m and 850\,$\mu$m maps. Fig. \ref{Fig:SCUBA2_850_CO} shows the central 4.5\arcmin$\times$4.5\arcmin\ field extracted from the 850-$\mu$m SCUBA2 map. The SCUBA2 maps reveal a dense core associated with IRAS~18144. We measure an FWHM of 18.5\arcsec\ and 22.6\arcsec\ respectively at 450\,$\mu$m and 850\,$\mu$m. These are larger than the average FWHM of 9.35\arcsec\ and 14.2\arcsec\ respectively at 450\,$\mu$m and 850\,$\mu$m for two point source standards observed prior to the target observation. The centroid of the core is at $\alpha$=18:17:24.01, $\delta$=-17:22:12.05. This location is only 3.4\arcsec\ from source `B' and is closer to source `B' than to `A'. Also plotted on the figure are the contours of the red- and blue-shifted lobes of the CO outflow (see \S\,\ref{COdata}), which appear to be centered on the sub-mm source. The maps also reveal filamentary emission extending in the NE and SW of the central source, more prominent towards the NE. The longest filament seen in our maps extends to a distance of 2.5\arcmin\ from the central source, which is 3.15\,pc. Table \ref{tab:SCUBA2_flx} shows the flux density measured in different circular apertures around the point source. We measure a peak flux density of 31\,Jy\,beam$^{-1}$ and 4\,Jy\,beam$^{-1}$ at 450 and 850\,$\mu$m respectively on the point source. \begin{figure} \centering \includegraphics[width=8.4cm,clip]{IRAS18144_850_CO.pdf} \caption{A 4.5\arcmin$\times$4.5\arcmin field extracted from the SCUBA2 850$\mu$m map. The dashed black contours show flux densities of 12, 3, 1, 0.5 and 0.1\,mJy~arcsec$^{-2}$ respectively on the 850$\mu$m map. The blue and red contours are generated from the CO(3--2) intensity maps of the blue- and red-shifted lobes of the outflow detected in our JCMT+HARP observations, integrated in velocity ranges of 23.3 -- 44.0\,km~s$^{-1}$ and 50.6 -- 71.3\,km~s$^{-1}$ respectively. The blue contours are at 51, 40 and 25\,K\,km\,s$^{-1}$ and the red contours are at 34.5, 30 and 15\,K\,km\,s$^{-1}$. A 1\arcmin$\times$1\arcmin image of the central source is shown in high contrast in the inset. `$\ast$' shows the location of source `A' and `$\times$' shows the location of source `B'.} \label{Fig:SCUBA2_850_CO} \end{figure} \begin{table} \caption{Flux density of `A+B' measured in different circular apertures from our SCUBA2 maps} \label{tab:SCUBA2_flx} \begin{tabular}{@{}lll@{}} \hline Aperture diameter & \multicolumn{2}{c}Flux density (Jy)\\ (arcsec) & 450\,$\mu$m & 850\,$\mu$m \\ \hline 30 & 66.6 & 5.8 \\ 40 & 82.2 & 7.4 \\ 45 & 87.9 & 8.0 \\ 60 & 99.6 & 9.3 \\ 80 & 109.4 & 10.4 \\ 160 & 120.1 & 11.8 \\ \hline \end{tabular} \end{table} \subsection{Heterodyne observations using HARP} \subsubsection {CO\,($J$=3--2) data} \label{COdata} We obtained heterodyne mapping observations in CO with the JCMT on 20140731~UT. The observations were performed using HARP \citep{buckle09} as the front end and the ASCIS autocorrelator as the back end, in position-switched raster-scan mode with quarter array spacing. $^{12}$CO\,(3--2) (345.796\,GHz) and H$^{13}$CN (345.3398\,GHz), and $^{13}$CO\,(3--2) (330.588\,GHz) and C$^{18}$O~(3--2) (329.331\,GHz) were simultaneously observed in two different dual sub-band settings of ASCIS. This gives a bandwidth of 250\,MHz and a resolution of 61\,kHz (0.055\,km~s$^{-1}$) for each line. The scans were performed in a 3$\arcmin\times$3$\arcmin$ area. The CSO $\tau_{225\,GHz}$ was 0.12 for $^{12}$CO/H$^{13}$CN observations, and 0.128 for $^{13}$CO/C$^{18}$O observations. The pointing accuracy on source {is better than 2\arcsec.} The data were reduced using the {\sc oracdr} pipeline, which performs a quality-assurance check on the time-series data, followed by trimming and de-spiking, and an iterative baseline removal routine before creating the final group files. After binning to a resolution of 0.423\,km~s$^{-1}$ (please see the next paragraph), we achieve an RMS noise level of 0.32, 0.31, 0.81 and 1.27\,K in $^{12}$CO, H$^{13}$CN, $^{13}$CO and C$^{18}$O respectively. We also downloaded JCMT archival data in $^{12}$CO\,(3--2), $^{13}$CO\,(3--2) and C$^{18}$O~(3--2) lines. The $^{12}$CO observations were performed on 20080419~UT with a 1000\,MHz band width 2048 channel set up giving a velocity resolution of 0.423\,km~s$^{-1}$. The CSO $\tau_{225\,GHz}$ was 0.115. The $^{13}$CO and C$^{18}$O lines were observed on 20080706~UT in a single setting with 250\,MHz band width 4096 channels for each line, giving a velocity resolution of 0.055 and 0.056\,km~s$^{-1}$ respectively. Similar to our observations, the archival observations were also obtained in position-switched raster scan mode, albeit with half array spacing. The CSO $\tau_{225\,GHz}$ was 0.09 for the $^{13}$CO/C$^{18}$O observations. These spectra were binned to a 0.423\,km~s$^{-1}$ resolution as in the $^{12}$CO line. The $^{12}$CO, $^{13}$CO and C$^{18}$O archival spectra reached an RMS noise level of 0.41, 0.56 and 0.88\,K respectively. We used our new $^{12}$CO observations for the figures presented in this paper. Column densities were estimated separately from our own observations and the archival observations to show the repeatability, and the average was used in the discussions. The results from the analysis of the CO data are presented in \S\ref{CO-outflow}. No H$^{13}$CN line emission was detected in our observation. \subsubsection{HCO$^{+}$\,(4-3)} \label{HCO+} Archival data of IRAS~18144 in the HCO$^+$\,(4-3) line (356.7\,GHz) were downloaded from the JCMT science archive at CADC. The observations were performed on 20100405\,UT with HARP in jiggle map mode covering a field of 2\arcmin$\times$2\arcmin. A 256\,MHz 4096 channel HARP setup was used giving a velocity resolution of 0.051\,km~s$^{-1}$. The mean atmospheric opacity at 225\,GHz, measured with the CSO dipper, was 0.084 for the period of the observations. The archival data was reduced using {\sc{oracdr}}, further reduction was carried out using {\sc{kappa}}. After binning the velocity resolution down to 0.423\,km~s$^{-1}$, we get an RMS of 0.11\,K. Fig. \ref{Fig_IRAS18144_HCO} shows the spectrum (at the peak) after binning. \begin{figure} \centering \includegraphics[width=8.4cm,clip]{IRAS18144_HCO_JCMT20100405.pdf} \caption{The HCO$^{+}$ line at the peak, binned down to a velocity resolution of 0.42~km~s$^{-1}$. The vertical line shows the systemic velocity of 47.3~km~s$^{-1}$. The peak T$_{mb}$ observed was 5.17\,K (T$_A^*$=3.31)} \label{Fig_IRAS18144_HCO} \end{figure} | \begin{enumerate} \item Our sub-mm observations of IRAS~18144 reveal massive star formation taking place in a dense core located in an isolated molecular cloud. \item We discover an E--W outflow in the CO(3--2) line. The CO outflow is aligned with the jet imaged by us in the 2.122\,$\mu$m H$_2$ line from this region showing that the outflow is jet-driven. We estimate a lower limit of 5.8\,M$_\odot$ for the mass in the outflow. \item Analysis of the near-IR spectrum of the jet discovered in H$_2$ gives and excitation temperature of 2440\,K and a foreground extinction of A$_V$=15.6 mag for the H$_2$ line emission lobe. \item Through near- and mid-IR imaging, we discover that IRAS~18144 hosts at least two massive YSOs (labelled `A' and `B') in a dense core mapped at 450 and 850\,$\mu$m. The newly discovered embedded source `B' is younger and is likely to be the one driving the outflow seen in CO and H$_2$. Both sources appear to be accreting, and are likely to be in class II and class I stages respectively. \item The $HK$ spectrum of source `A' shows a steeply rising SED, with Br$_\gamma$ emission, which is typical of a YSO with accretion. Assuming that the Br$_\gamma$ emission is from the accretion disk, we calculate an accretion rate of 2.2$\times$10$^{-5}$\,M$_\odot$~year$^{-1}$. \item Source `B' is resolved well from `A' in our mid-IR observations. The spatial resolution of the far-IR and sub-mm data are not sufficient to resolve the two sources. Assuming that `B' is the sole contributor to the emission observed at far-IR and sub-mm wavelengths, we modelled the SED, which revealed a 16.2\,M$_\odot$ YSO with a luminosity of 3.07$\times$10$^4$\,L$_\odot$. The disk parameters exhibited from the SED analysis have large uncertainties. It is likely that our assumption that `B' is the sole contributor to the observed flux densities at longer wavelengths is not justified, and that the contribution from source `A', at those wavelengths cannot be neglected. Even though the two sources appear close at 2.6\arcsec, they have a projected separation of 11258\,AU at a distance of 4.33\,kpc. We also cannot ignore the possiblity of `B' being composed of more than one source, which are not resolved in our observations. High angular resolution observations at far-IR and sub-mm wavelengths are required to reliably estimate the parameters of the YSOs, and to learn if there are more YSOs associated with this cloud. \item An examination of the near-IR magnitudes available for source `A', spanning a period of $\sim$20 years, show that source brightened by over 1.1 magnitudes and is now fading. The variation of $H-K$ suggests that the variability is due to varying amount of obscuration by circumstellar dust. \item We conclude that IRAS~18144 hosts at least two massive YSOs, located inside a dense core situated in an isolated cloud. We estimate a virial mass of 1007\,M$_{\odot}$ for the core from archival HCO$^+$(4--3) data. This is comparable to a mass of 883\,M$_{\odot}$ derived from our SCUBA-2 850\,$\mu$m observations. The SCUBA-2 continuum maps and HCO$^+$(4--3) maps show the densest region of the cloud (core) where the two YSOs are located. From our $^{12}$CO(3-2) data, we estimate a lower limit of 3994\,M$_\odot$ for the mass of the cloud hosting the core. \end{enumerate} | 18 | 8 | 1808.00230 |
1808 | 1808.07626_arXiv.txt | The infrared signatures of polycyclic aromatic hydrocarbons (PAHs) are observed in a variety of astrophysical objects, including the circumnuclear medium of active galactic nuclei (AGNs). These are sources of highly energetic photons (0.2 to 10 keV), exposing the PAHs to a harsh environment. In this work, we examined experimentally the photoionization and photostability of naphthalene (C$_{10}$H$_{8}$), anthracene (C$_{14}$H$_{10}$), 2-methyl-anthracene (C$_{14}$H$_{9}$CH$_{3}$) and pyrene (C$_{16}$H$_{10}$) upon interaction with photons of 275, 310 and 2500 eV. The measurements were performed at the Brazilian Synchrotron Light Laboratory using time-of-flight mass-spectrometry (TOF-MS). We determined the absolute photoionization and photodissociation cross sections as a function of the incident photon energy; the production rates of singly, doubly and triply charged ions; and the molecular half-lives in regions surrounding AGNs. Even considering moderate X-ray optical depth values ($\tau = 4.45$) due to attenuation by the dusty torus, the half-lives are not long enough to account for PAH detection. Our results suggest that a more sophisticated interplay between PAHs and dust grains should be present in order to circumvent molecular destruction. We could not see any significant difference in the half-life values by increasing the size of the PAH carbon backbone, N$_C$, from 10 to 16. However, we show that the multiple photoionization rates are significantly greater than the single ones, irrespective of the AGN source. We suggest that an enrichment of multiply charged ions caused by X-rays can occur in AGNs. | Neutral and ionized Polycyclic Aromatic Hydrocarbons (PAHs) are detected in astronomical sources through emission bands in the infrared (IR) wavelength range, due to the corresponding molecular vibrations. \citep{Peeters2002a,Tielens2008}. The main vibrational features of PAHs comprise the \-{C$-$H} (3.3 $\mu$m), \-{C$=$C} (6.2 $\mu$m) and \-{C$-$C} (7.7 $\mu$m) bond stretching modes, as well as the in-plane (8.6 $\mu$m) and out-of-plane (11.3 $\mu$m) \-{C$-$H} bond bending modes. From the analysis of such IR bands, PAHs have been observed in a diversity of galactic objects, such as planetary nebulae \citep{Waters1998,Gorny2001,Ohsawa2012,Guzman-Ramirez2014}, HII regions \citep{Roelfsema1996,Peeters2002}, reflection nebulae \citep{Boersma2014,Ricca2018}, protoplanetary disks \citep{Ressler2003,Maaskant2014,Schworer2017,Seok2017,Taha2018}, among others. The PAH IR features contribute to about 10\% of the ISM luminosity in the 1-1000 $\mu$m range, accounting for a large fraction of the elemental C in star-forming galaxies \citep{Lagache2004}. In addition, their luminosities are well correlated with star formation rates \citep{Peeters2004,Stierwalt2014,Alonso-Herrero2014,Esparza-Arredondo2018}. In view of such remarkable characteristics, continuing research on the formation and stability of PAHs has a key relevance to astrochemistry. In addition to galactic sources, PAHs have also been observed in a variety of extragalactic objects, such as HII regions in the Magellanic Clouds \citep{Li2002a,Vermeij2002,Oey2017}, local dusty elliptical galaxies \citep{Kaneda2008}, starburst galaxies \citep{Brandl2006}, submillimeter galaxies (SMGs) \citep{Menendez-Delmestre2009}, ultra-luminous IR galaxies (ULIRGs) \citep{Desai2007} and in the circumnuclear regions of Active Galactic Nuclei (AGNs) \citep{Lutz1998,ODowd2009,Tommasin2010,Sales2013,Esquej2013,Alonso-Herrero2016}. Concerning the latter objects, PAHs have been identified in both Seyfert 1 and Seyfert 2 galaxies \citep{Mazzarella1994,Deo2007, Diamond-Stanic2010, Sales2013}, in low-ionization nuclear emission-line regions (LINERs) \citep{Sturm2006} and obscured quasars \citep{MartinezSansigre2008}. AGNs are sources of both soft (0.2-2 keV) and tender (2-10 keV) X-ray radiation, and contribute significantly to the extragalactic X-ray background \citep{Comastri1994,Mushotzky2000,Lubinski2016,Hickox2018}. The X-ray emission mechanism is known to be powered by gas accretion on to a central supermassive black hole \citep{Ferrarese2005,Zhang2005,DiMatteo2005,Barai2012}, which is the basis of the standard unification model for AGNs \citep{Antonucci1993,Urry1995} --- for a recent review, see \citealt{Netzer2015}. The combination of a 0.2-10 keV radiation field and the presence of PAHs in the AGN vicinity provides an interesting scenario in which laboratory investigation of molecular photoionization and photodestruction in the X-ray energy range could provide important information. The experimental studies of vacuum ultraviolet (VUV) photoionization and photodissociation of PAHs were firstly reported at about thirty years ago by \cite{Leach1989, Leach1989b}, using time-of-flight mass-spectrometry and photoelectron-photoion coincidence techniques. \cite{Jochims1994,Jochims1996,Jochims1999} have also investigated the VUV photostability of PAHs and methyl-substituted PAHs. In addition, the competition between VUV photoionization, photofragmentation and the photoproduction of dications in the context of interstellar PAH population were discussed by \citealt{Zhen2016}. Using soft X-ray photons with energies around the C1s$\rightarrow \pi$* resonance (285 eV), \cite{Reitsma2014,Reitsma2015} have investigated the fragmentation of PAHs cations. However, from our perspective, there is still a lack of information on the ionization and dissociation of PAHs driven by photons of higher energies, which would be useful in the context of the circumnuclear regions of AGNs. The photochemistry of these regions should be highly affected by a radiation field with 0.2-10 keV photons due to its high penetration power, even through a gas with column densities up to 10$^{24}$ cm$^{-2}$, such as X-ray dominated regions (XRD, see \citealt{Usero2004}) . \begin{figure} \centering \includegraphics[width=0.7\columnwidth]{fig1.eps} \caption{The polycyclic aromatic hydrocarbons studied in this work.} \label{fig:1molecules} \end{figure} The present work aims at studying the effect of X-ray photon interactions on the photodissociation and photoionization rates of the PAH molecules shown in Fig. \ref{fig:1molecules}. In order to compare the ionic fragmentation profiles with respect to the photon energy, we have measured mass spectra at selected energies below (275 eV) and above (310 eV) the C1s resonance features of these PAHs, and at 2500 eV. From the mass spectra obtained herein, we determined the photoionization and photodissociation cross-sections. These results are discussed in the context of the chemistry of the circumnuclear regions of AGNs, which is the focus of our study. | \label{sec:summary} In this work, we examined the photoionization and photodissociation profiles of selected polycyclic aromatic hydrocarbons (PAHs) upon their interaction with soft and tender X-rays. For this purpose, mass spectra of naphthalene (C$_{10}$H$_{8}$), anthracene (C$_{14}$H$_{10}$), 2-methyl-anthracene (C$_{14}$H$_{9}$CH$_{3}$, or C$_{15}$H$_{12}$) and pyrene (C$_{16}$H$_{10}$) were obtained for energies below (275 eV) and above (310 eV) the C1s resonance features of these PAHs, and at 2500 eV. The results are discussed in the context of the chemistry of the circumnuclear regions of AGNs, which is the focus of our study. The measurements were performed at the Brazilian Synchrotron Light Laboratory (LNLS) using time-of-flight mass-spectrometry and photoelectron-photoion coincidence techniques. Two processes emerge as the main channels after excitation and photoionization of the molecules: the breakage of the carbon backbone, in which singly charged ions with m/z lower than half of the parent molecule are mostly produced, and the formation of multiply charged ions, for which the initial PAH carbon content is predominantly preserved. Multiple non-dissociative photoionization is more pronouncedly activated if the photon energy is increased to the keV region. The same trend is observed if the carbon backbone size of the molecule is expanded. In fact, multiple ionization seems to be significantly more efficient upon the interaction with 2500 eV photons than by collision with high energy protons and $\alpha$-particles. On the other hand, methylation of the PAH does not seem to particularly affect the formation of multiply charged species in the keV region in comparison to its non-methylated analog. The distribution of multiply charged species that are produced upon interaction with a 2.5 keV photon seems to be more dependent on their intrinsic stabilities, rather than on the characteristics of the PAH parent molecule. Taken together, these results could indicate that high energy photons are able to trigger the formation of multiply charged PAHs in the surroundings of AGNs. By taking our $PIY$ results, we could determine the photoionization and photodissociation cross sections of the PAH molecules at 2500 eV. These values were used to estimate the photoionization and photodissociation rates of PAHs in the circumnuclear regions (20-80 pc) of six AGN sources with distinct X-ray fluxes. From the photodissociation rates, we could estimate the half-lives of those molecules for different optical depth values of the X-ray photon flux. These values were compared to the PAH injection timescale ($2.5 \times 10^9$ yr) described by \cite{Jones1994} assuming that the main sources of PAHs are carbon-rich AGB stars. In spite of considering attenuation of the X-ray radiation field by a dusty torus associated with an H$_2$ column density of $2 \times 10^{23}$ cm$^{-2}$ ($\tau = 4.45$), the lifetime of PAHs spanned values from 10$^8$ yr to merely 10$^2$ yr. These results may indicate that, in order to circumvent molecular destruction, a more sophisticated interplay between PAHs and dust grains should be considered. In this perspective, we briefly describe two possible scenarios in which grains could assist in the survival of PAHs. We could not see any significant difference in the half-life values by increasing the size of the carbon backbone. This is probably due to the fact that we spanned only small-size PAH molecules ($10 \leq \text{N}_C \leq 16$). In addition, we show that the multiple photoionization rates are significantly greater than the single ones, irrespective of the AGN source. These results suggest that an enrichment of multiply charged ions caused by X-ray photoselection can occur in AGNs. The precise determination of the charge state of PAHs based on specific spectral signatures should be developed in order to confirm this photoselectivity mechanism. | 18 | 8 | 1808.07626 |
1808 | 1808.03280_arXiv.txt | We report a Giant Metrewave Radio Telescope (GMRT) survey for associated H{\sc i} 21-cm absorption from 50 active galactic nuclei (AGNs), at $z \approx 0.04 - 3.01$, selected from the Caltech-Jodrell Bank Flat-spectrum (CJF) sample. Clean spectra were obtained towards 40 sources, yielding two new absorption detections, at $z = 0.229$ towards TXS~0003+380 and $z = 0.333$ towards TXS~1456+375, besides confirming an earlier detection, at $z = 1.277$ towards TXS~1543+480. There are 92 CJF sources, at $0.01 \lesssim z \lesssim 3.6$, with searches for associated H{\sc i} 21-cm absorption, by far the largest uniformly-selected AGN sample with searches for such absorption. We find weak ($\approx 2\sigma$) evidence for a lower detection rate of H{\sc i} 21-cm absorption at high redshifts, with detection rates of $28^{+10}_{-8}$\% and $7^{+6}_{-4}$\% in the low-$z$ ($z < z_{\rm med}$) and high-$z$ ($z > z_{\rm med}$) sub-samples, respectively. We use two-sample tests to find that the strength of the H{\sc i} 21-cm absorption in the AGNs of our sample depends on both redshift and AGN luminosity, with a lower detection rate and weaker absorption at high redshifts and high ultraviolet/radio AGN luminosities. Unfortunately, the luminosity bias in our sample, with high-luminosity AGNs arising at high redshifts, implies that it is not currently possible to identify whether redshift evolution or AGN luminosity is the primary cause of the weaker absorption in high-$z$, high-luminosity AGNs. We find that the strength of H{\sc i} 21-cm absorption does not depend on AGN colour, suggesting that dust extinction is not the main cause of reddening in the CJF sample. | Neutral hydrogen (\hi) is an important constituent of the gas in the environments of active galactic nuclei (AGNs). For radio-loud AGNs, the \hii\ transition allows the presence of such gas to be discerned, and its kinematics studied, via \hii\ absorption studies against the AGN radio continuum. Such ``associated'' \hii\ absorption studies allow a detailed probe of physical conditions in AGN environments \citep[see, e.g., ][for a recent review]{morganti18}, and their evolution with redshift. For example, one can test whether the strength of the \hii\ absorption depends on the nature of the AGN (e.g. core-dominated or lobe-dominated, high-luminosity or low-luminosity, flat spectrum or steep spectrum, etc), providing information on the distribution and the excitation of neutral gas in different AGN environments. One might also test whether the strength of \hii\ absorption depends on redshift: weaker \hii\ absorption in typical AGNs over some range of redshifts might indicate a paucity of gas in the AGN environments at this epoch, shedding light on the gas accretion process and the fuelling of AGNs. The \hii\ absorption velocity, relative to the AGN redshift, also contains information on local conditions, with redshifted (relative to the AGN) absorption indicating the presence of inflowing gas, and blueshifted absorption, a signature of gas outflows. If the associated \hii\ absorption systematically arises at higher velocities than the AGN redshift, it would suggest that neutral gas is predominantly flowing towards typical AGNs at this redshift, which could fuel the nuclear activity \citep[e.g.][]{vangorkom89}. Conversely, mostly blueshifted \hii\ absorption would imply a predominance of gas outflows, that might result in shutting down of the AGN activity \citep[e.g.][]{vermeulen03,gupta06}. Similarly, narrow \hii\ lines with velocity spreads $\lesssim 100$ km~s$^{-1}$ and relatively small velocity offsets from the AGN systemic redshift are likely to indicate gas clouds with low velocity dispersion, perhaps rotating in a circumnuclear disk \citep[e.g.][]{dwaraka95,conway99,gereb15}. Broad absorption features, with widths $\approx 200-300$~\kms\ suggest the presence of unsettled gas, possibly interacting with the active nucleus. Finally, the broadest lines, with widths $\gtrsim 500$~\kms, are likely to arise due to interactions between the neutral gas and the radio jets in powerful radio sources \citep[e.g.][]{morganti03,morganti05,mahony13}. A further important advantage of \hii\ absorption studies in probing AGN environments stems from the fact that the radio emission is often extended, stemming from both the AGN core and structures such as jets, lobes, hotspots, etc. This allows the exciting possibility of using very long baseline interferometry (VLBI) techniques to map the \hii\ absorption against the extended radio continuum, on scales of $\approx 10-1000$~pc, to determine the spatial structure of the absorbing gas, and connections with the radio continuum \citep[e.g.][]{mundell95,carilli98c,peck98,peck99,conway99,beswick02,morganti04,labiano06,struve12,morganti13}. Such VLBI absorption mapping studies are critical to identify the location of the absorbing gas in the AGN environment. Since the first detection of associated \hii\ absorption, by \citet{roberts70} in NGC5128, a large number of searches for associated \hii\ absorption have been carried out, using a variety of radio telescopes. More than 400 AGNs have so far been searched for associated \hii\ absorption, with more than seventy-five detections \citep[e.g.][]{dickey86,vangorkom89,carilli98,moore99,gallimore99,morganti01,morganti05,pihlstrom03,vermeulen03,gupta06,chandola11,allison12,gereb15,maccagni17,aditya18}. Molecular absorption has also been detected in a handful of these systems, indicating the presence of dense gas \citep[e.g.][]{gardner76,wiklind94,wiklind96,kanekar02,kanekar08c}. However, the vast majority of the detections of \hii\ absorption are at low redshifts, $z < 0.25$ \citep[e.g.][]{vangorkom89,gereb15,maccagni17}: the current sample of associated \hii\ absorbers is dominated by the 59 detections of \citet{gereb15} and \citet{maccagni17}, at $0.05 < z < 0.25$. Further, most searches for redshifted \hii\ absorption have been in highly heterogeneous samples, making it hard to differentiate between the causes for the presence or absence of absorption (e.g. redshift, AGN type, AGN luminosity, etc.). We have hence been using the Giant Metrewave Radio Telescope (GMRT) survey for redshifted associated \hii\ absorption towards a uniformly-selected AGN sample, distributed over a wide redshift range. Earlier results from this survey were reported in \citet{aditya16} and \citet{aditya17}. | \label{sec:discuss} \subsection{A uniformly-selected flat-spectrum sample} \label{sec:sample} Searches for redshifted associated \hii\ absorption have so far been carried out in more than 400 AGNs, with $\gtrsim 75$ detections of \hii\ absorption \citep[e.g.][]{vangorkom89,vermeulen03,gupta06,curran10,gereb15,maccagni17}. The vast majority of both searches and detections are at low redshifts, $z < 1$, where the typical detection fraction is $\approx 30$\% \citep[e.g.][]{pihlstrom03,gupta06,maccagni17}. The situation is very different at high redshifts, $z > 1$, with searches for associated \hii\ absorption in only $\approx 25$~AGNs \citep[e.g.][]{gupta06,curran13} and just four detections in the literature, prior to our survey. The implied detection rate of \hii\ absorption is $\approx 16^{+13}_{-8}$\%, where the errors are from Poisson statistics \citep[][]{gehrels86}. While the detection rate of \hii\ absorption at $z > 1$ is only half that at low redshifts, the difference between the two is not statistically significant due to the large uncertainty in the high-$z$ value, simply due to the fact that few high-$z$ AGNs have hitherto been targetted in \hii\ absorption studies. An important part of the present survey was simply to increase the number of searches for redshifted \hii\ absorption at $z > 1$. If high-$z$ AGNs yield a detection fraction similar to those at low redshifts, the survey would then yield a large sample of associated \hii\ absorbers, suitable for detailed kinematic studies of AGN environments. Conversely, if the detection fraction remains low at $z > 1$, this would be evidence for redshift evolution in AGN environments. Further, most studies of associated \hii\ absorption \citep[e.g.][]{vermeulen03,gupta06,curran13,maccagni17} have targetted highly heterogeneous AGN samples at all redshifts. This heterogeneity makes it difficult to distinguish between possible redshift evolution in the AGN environment and differences in the AGN samples at different redshifts. For example, \citet[][]{gupta06} carried out an analysis of 96 AGNs, mostly at low to intermediate redshifts, $z < 1$. They found little evidence for redshift evolution in either the detection rates of \hii\ absorption or in the distribution of the \hii\ optical depths. However, their sample included a range of AGN types, with 21 large radio galaxies, 13 flat-spectrum sources, 35 CSS sources, and 27 GPS sources. The heterogeneity of the sample makes it difficult to reliably interpret the observational data. An alternative explanation for the tentative result that the strength of associated \hii\ absorption may be weaker at high redshifts stems from the luminosity bias in most AGN samples used in such studies. High-$z$ AGN samples typically contain more objects with higher rest-frame UV and radio luminosities. \citet[][]{curran08} suggest that the high AGN luminosity in the UV and/or radio wavebands could lead to a lower \hii\ optical depth, either by ionizing the \hi\ (and thus reducing the \hi\ column density) or by altering the hyperfine level populations (and thus increasing the spin temperature). \citet[][]{curran08} hence argued that the high luminosities of high-$z$ AGNs may be the primary cause of the low detection rate of associated \hii\ absorption at high redshifts \citep[see also][]{curran13}. However, we note that the AGN samples of \citet{curran08} and \citet{curran13} were also highly heterogeneous, containing all the AGNs that had been searched for \hii\ absorption in the literature. The present survey for \hii\ absorption aims to address the above issues by carrying out a search for \hii\ absorption in a large and uniformly-selected sample of AGNs, selected from the CJF catalogue, to investigate the dependence of the \hii\ absorption strength and detectability on redshift and AGN properties (e.g. luminosity, colour, spectral index, etc). In our pilot study \citep[][]{aditya16}, we combined 23 new searches for redshifted \hii\ absorption with 29 searches from the literature (all in AGNs from the CJF sample), to find tentative evidence (at $\approx 3\sigma$ significance) that the strength of associated \hii\ absorption depends on both redshift and AGN luminosity, with weaker \hii\ absorption obtained at both high redshifts and higher AGN luminosities. In the present work, we have completed our GMRT \hii\ absorption survey of AGNs of the CJF sample, targetting nearly all CJF sources whose redshifted \hii\ line frequencies lie within the GMRT's legacy 327, 610, and 1420-MHz bands. Including 39 sources with usable data from this study, our full sample consists of 92 flat-spectrum AGNs, which includes 63 sources from our GMRT observations and 29 from the literature. The sample contains 16 detections of \hii\ absorption [including the three confirmed detections of this paper towards TXS~0003+380, TXS~1456+375, and TXS~1543+480, the confirmed detection towards TXS~1954+513 \citep{aditya17}, and our tentative detection towards TXS~0604+728 \citep{aditya16}], and 76 non-detections, yielding upper limits to the \hii\ optical depth. This is by far the largest sample of uniformly-selected AGNs that have been searched for associated \hii\ absorption. The sample covers a large redshift range, $0.01 \lesssim z \lesssim 3.6$, with more than half of the sample ($\approx 50$ sources) at $z > 1$. The 92 CJF sources of our full sample are listed in order of increasing redshift, in Table~\ref{table:cjf}. The columns of this table are (1)~the AGN name, (2)~the AGN redshift, (3)~the velocity-integrated \hii\ optical depth in \kms, or, for non-detections of \hii\ absorption, the 3$\sigma$ upper limit to the \hii\ optical depth, assuming a Gaussian profile with a line FWHM of 100~\kms, (4)~the rest-frame 1216~\AA\ AGN luminosity $L'_{UV} \equiv Log[L_{UV}/(W\;Hz^{-1})]$, inferred by interpolating between measured luminosities in UV and/or optical wavebands (see below), (5)~the rest-frame 1.4~GHz AGN luminosity $L'_{\rm 1.4\;GHz} \equiv {\rm Log}[L_{\rm 1.4\;GHz}/({\rm W\;Hz^{-1}})]$, (6)~the AGN spectral index close to the redshifted \hii\ line frequency, $\rm \alpha_{\rm 21-cm}$, (7)~the AGN colour (R$-$K) between the R- and K-bands, (8)~the reference for the search for \hii\ absorption, and (9)~references for the estimates of UV, optical, and near-infrared (NIR) luminosities, which were used to infer $L'_{\rm UV}$ and (R$-$K). The rest-frame 1216~\AA\ AGN luminosity was estimated following the procedure of \citet[][]{curran10}. We first determined the flux density $F_{\rm UV}$ at $1216\times(1 + z)$~\AA\ for each AGN, using a power-law spectrum to interpolate between measured flux densities at two nearby optical and/or UV wavebands from the literature. The luminosity at rest-frame 1216~\AA\ was then inferred from the expression ${L_{UV}} = {4\pi D_{\rm AGN}^{2} F_{\rm UV}}/(1+z)$, where $D_{\rm AGN}$ is the AGN's luminosity distance. For two AGNs, TXS~0424+670 and TXS~1020+400, the flux density is known only at a single optical waveband, quite distant from the redshifted 1216~\AA\ wavelength. These systems hence do not have a listed rest-frame 1216~\AA\ luminosity in Table~\ref{table:cjf}. The radio spectral indices of the AGNs were computed from their flux densities at the redshifted \hii\ line frequency and a nearby frequency at which a flux density estimate was available in the literature. For all sources at $z > 1$, the second frequency was 1.4~GHz, from the FIRST or NVSS surveys \citep{becker95,condon98}. For sources at $z < 1$, the second frequency was either 365~MHz \citep[the Texas survey;][]{douglas96} or 325~MHz \citep[the WENSS survey;][]{rengelink97}. Finally, the (R$-$K) colour could only be inferred for 58 AGNs of the full sample; the remaining sources do not have K-band information in the literature. In the following sections, we will examine the dependence of the \hii\ detection fraction and the distribution of integrated \hii\ optical depth on redshift, radio spectral index, AGN radio and UV luminosities, and the (R$-$K) colour, for the full sample of 92 flat-spectrum sources. \setcounter{table}{1} \begin{table*} \footnotesize \caption{The 92 CJF sources of the full sample with searches for redshifted \hii\ absorption, listed in order of increasing redshift. 63 sources are from the present survey, 39 from this paper and 24 from \citet{aditya16,aditya17}, while 29 are from the literature. See main text for discussion. \label{table:cjf}} \begin{center} \begin{tabular}{|lcccccccc|} \hline \\ AGN & $z$ & $\int \tau dv$ & $L'_{\rm UV}$$^a$~\tnote{a} & $L'_{\rm 1.4\;GHz}$ & $\alpha_{\rm 21-cm}$ & (R$-$K)$^b$~\tnote{b} & Refs.$^c$~\tnote{c} & Refs.$^d$~\tnote{d} \\ & & \kms & & & & & (21-cm) & (UV,Opt.,NIR) \\ \hline TXS 1146+596 & 0.011 & $5.3 \pm 1.8$ & 20.30 & 23.09 & 0.26 & 2.3 & 3 & 1--4 \\ TXS 0316+413 & 0.018 & $1.3$ & 20.61 & 25.13 & 1.05 & -0.2 & 4 & 1,3,4 \\ B3 0651+410 & 0.022 & $<0.82$ & 18.69 & 23.35 & 0.60 & -1.0 & 5 & 1,3,4 \\ TXS 1101+384 & 0.030 & $<0.63$ & 21.28 & 24.18 & 0.17 & -1.7 & 6 & 4,5 \\ TXS 1744+557 & 0.030 & $<1.2$ & 19.53 & 24.02 & 0.22 & -3.1 & 7 & 1,3,4 \\ TXS 1652+398 & 0.034 & $<2.4$ & 21.58 & 24.56 & -0.12 & -1.3 & 6 & 1,3,4 \\ TXS 0344+405 & 0.039 & $<6.3$ & 21.34 & 23.23 & 1.70 & 0.6 & 1 & 1,3,4 \\ TXS 0733+597 & 0.041 & $<0.67$ & 19.69 & 24.24 & -0.28 & -1.8 & 1 & 1,3,4 \\ TXS 1254+571 & 0.042 & $33.9 \pm 3.8$ & 20.51 & 24.06 & -0.57 & 6.1 & 14 & 2--4 \\ TXS 1807+698 & 0.051 & $<$1.6 & 20.91 & 25.00 & 1.05 & 0.2 & 6 & 1,3,4 \\ TXS 0402+379 & 0.055 & $0.98 \pm 0.11$ & 19.88 & 25.31 & -0.29 & 3.6 & 9 & 3,6 \\ TXS 1144+352 & 0.063 & $<1.1$ & 20.26 & 24.75 & 0.54 & -0.9 & 7 & 1,3,4 \\ TXS 2200+420 & 0.069 & $<0.95$ & 20.30 & 25.78 & -0.14 & 2.0 & 6 & 7--10 \\ S5 2116+81 & 0.084 & $<1.4$ & 21.25 & 24.31 & -0.67 & 1.5 & 1 & 1,3,4 \\ TXS 1418+546 & 0.153 & $<1.5$ & 21.71 & 25.48 & 0.58 & 4.7 & 1 & 1,4,11 \\ TXS 1946+708 & 0.101 & $15.8 \pm 4.6$ & 18.60 & 25.35 & -0.33 & 1.3 & 10 & 3,4 \\ TXS 0309+411 & 0.134 & $<0.92$ & 18.48 & 25.23 & 0.08 & 1.5 & 7 & 3,4 \\ S4 0749+54 & 0.200 & $<1.2$ & 20.71 & 25.68 & 0.35 & 2.5 & 1 & 1,4,11 \\ IVS B1622+665 & 0.201 & $<1.7$ & 20.00 & 25.24 & 0.53 & 2.9 & 5 & 3,4 \\ S5 1826+79 & 0.224 & $<15$ & 21.28 & 25.57 & 0.59 & 2.1 & 11 & 1,3,4 \\ TXS 2021+614 & 0.227 & $<0.21$ & -- & 26.43 & 0.07 & 3.3 & 11 & 3,4 \\ TXS 0003+380 & 0.229 & $1.943 \pm 0.057$ & 20.62 & 25.80 & 0.49 & 3.6 & 1 & 1,3,4 \\ TXS 2352+495 & 0.238 & $1.7$ & 20.84 & 26.46 & -0.09 & 2.9 & 11 & 3,4,12 \\ TXS 0831+557 & 0.241 & $0.58$ & 21.39 & 27.03 & -0.15 & 0.3 & 11 & 1,3,4 \\ TXS 0010+405 & 0.255 & $<1.4$ & 21.24 & 25.72 & 1.37 & 2.2 & 1 & 1,3,4 \\ TXS 1719+357 & 0.263 & $<1.2$ & 21.90 & 25.61 & 0.74 & 3.2 & 1 & 1,3,4 \\ TXS 1943+546 & 0.263 & $2.9$ & 20.85 & 26.46 & -0.45 & 1.4 & 11 & 1,3,6 \\ TXS 0424+670 & 0.324 & $<0.89$ & -- & 26.24 & -0.46 & -- & 1 & -- \\ B3 1456+375 & 0.333 & $3.834 \pm 0.079$ & 20.65 & 25.58 & -0.41 & 5.1 & 1 & 1,3,7 \\ S5 2007+77 & 0.342 & $<9.9 $ & 22.03 & 26.25 & 0.09 & 3.5 & 1 & 3,7 \\ TXS 0035+367 & 0.366 & $<0.51$ & 21.56 & 26.26 & -0.17 & 2.4 & 1 & 1,3,4 \\ TXS 0954+658 & 0.368 & $<5.8 $ & 21.96 & 26.55 & 0.01 & 3.3 & 1 & 1,3,4 \\ CJ2 0925+504 & 0.370 & $<1.2 $ & 22.75 & 26.00 & 0.38 & 2.3 & 1 & 1,3,7 \\ TXS 0110+495 & 0.389 & $<0.43$ & 21.09 & 26.34 & -0.01 & 2.4 & 1 & 1,3,4 \\ TXS 1031+567 & 0.459 & $<0.76$ & 19.90 & 26.96 & -0.20 & -- & 11 & 2,13 \\ TXS 1355+441 & 0.646 & $19$ & 20.60 & 26.94 & -0.34 & -- & 11 & 1,2,13 \\ S4 0108+38 & 0.669 & $46 \pm 7$ & 21.63 & 26.95 & 1.16 & -- & 12 & 1,14,15 \\ TXS 1504+377 & 0.672 & $27.20 \pm 0.04$ & 20.30 & 26.95 & -0.21 & -- & 15 & 2 \\ TXS 0923+392 & 0.695 & $<0.54$ & 23.47 & 27.49 & -0.38 & 1.7 & 11 & 1,4,15 \\ TS5 0950+74 & 0.695 & $<1.4$ & 21.65 & 27.15 & 0.92 & -- & 11 & 3 \\ TXS 1642+690 & 0.751 & $<0.69$ & 22.66 & 27.34 & 0.03 & -- & 11 & 3 \\ S4 1843+35 & 0.764 & $<6.0$ & 24.66 & 27.55 & -0.03 & -- & 11 & 3 \\ TXS 1030+415 & 1.117 & $<0.69$ & 23.32 & 27.30 & -0.61 & -- & 1 & 2 \\ TXS 0600+442 & 1.136 & $<0.71$ & -- & 27.61 & -0.37 & -- & 2 & -- \\ S5 1044+71 & 1.150 & $<0.38$ & 23.26 & 27.70 & -0.97 & 4.7 & 1 & 1,3,4 \\ TXS 2356+390 & 1.198 & $<0.42$ & 22.36 & 27.36 & -1.16 & -- & 2 & 1,3 \\ TXS 0821+394 & 1.216 & $<0.44$ & 23.55 & 27.97 & -0.75 & 3.0 & 2 & 1--4 \\ TXS 1954+513 & 1.223 & $0.698 \pm 0.036$ & 23.23 & 27.75 & 0.06 & 2.4 & 16 & 3,4 \\ TXS 1105+437 & 1.226 & $<0.70$ & 23.04 & 27.18 & -0.38 & -- & 1 & 1,2 \\ TXS 1015+359 & 1.228 & $<0.43$ & 24.31 & 27.42 & -0.02 & -- & 1 & 1 \\ TXS 1432+422 & 1.240 & $<0.95$ & 22.62 & 27.04 & 0.10 & -- & 1 & 1,3,4 \\ TXS 0945+408 & 1.249 & $<0.30$ & 23.93 & 27.89 & -0.46 & 2.5 & 2 & 1,3,4 \\ S5 1150+81 & 1.250 & $<1.1 $ & 23.54 & 27.84 & -0.33 & 3.3 & 1 & 1,3,4 \\ TXS 1020+400 & 1.254 & $<0.51$ & -- & 27.77 & -0.78 & 3.1 & 1 & 3,4 \\ S5 1039+81 & 1.260 & $<1.5 $ & 23.90 & 27.46 & 0.01 & 2.9 & 1 & 3,4 \\ \hline \hline \end{tabular} \end{center} \end{table*} \setcounter{table}{1} \begin{table*} \caption{(contd.) \label{cjfcontd}} \begin{center} \begin{tabular}{|lcccccccc|} \hline \\ AGN & $z$ & $\int \tau dv$ & $L'_{\rm UV}$$^a$~\tnote{a} & $L'_{\rm 1.4\;GHz}$ & $\alpha_{\rm 21-cm}$ & (R$-$K)$^b$~\tnote{b} & Refs.$^c$~\tnote{c} & Refs.$^d$~\tnote{d} \\ & & \kms & & & & & (21-cm) & (UV,Opt.,NIR) \\ \hline TXS 0641+392 & 1.266 & $<0.68$ & 22.64 & 27.22 & 0.79 & -- & 2 & 3 \\ TXS 0537+531 & 1.275 & $<0.32$ & 22.84 & 27.46 & 0.01 & 2.7 & 2 & 1,3,4 \\ TXS 1543+480 & 1.277 & $9.69 \pm 0.53$ & 22.08 & 27.36 & -0.23 & 6.3 & 13 & 1,2,4 \\ TXS 1656+571 & 1.281 & $<0.74$ & 23.56 & 27.77 & -0.54 & -- & 1 & 1,3 \\ TXS 0707+476 & 1.292 & $<0.30$ & 23.86 & 27.63 & -0.15 & 1.4 & 2 & 1,3,4 \\ TXS 0833+416 & 1.301 & $<0.63$ & 24.06 & 27.26 & -0.02 & 0.2 & 1 & 1,3,4 \\ TXS 2319+444 & 1.310 & $<0.65$ & 22.41 & 27.20 & -0.05 & -- & 1 & 1,3 \\ TXS 0248+430 & 1.311 & $< 1.4$ & 23.10 & 27.84 & 0.22 & -0.9 & 3 & 1,3,4,16 \\ TXS 1240+381 & 1.318 & $<0.63$ & 23.53 & 27.37 & 0.01 & 2.9 & 1 & 1,3,4 \\ TXS 0850+581 & 1.318 & $<0.37$ & 23.53 & 27.63 & -0.26 & -- & 2 & 1 \\ TXS 2007+659 & 1.325 & $<0.76$ & 22.51 & 27.44 & -0.26 & -- & 1 & 1,3 \\ S5 2353+81 & 1.344 & $<0.87$ & 22.30 & 27.41 & -0.53 & 2.9 & 2 & 1,3,6 \\ JVAS J2236+7322 & 1.345 & $<0.97$ & 22.64 & 27.08 & -0.01 & -- & 1 & 1,3 \\ TXS 1342+663 & 1.351 & $<0.98$ & 22.74 & 27.02 & 1.23 & 4.3 & 1 & 1--4 \\ TXS 0035+413 & 1.353 & $<0.76$ & 23.12 & 27.39 & 0.35 & -- & 2 & 1 \\ TXS 1739+522 & 1.375 & $<0.58$ & 23.73 & 27.67 & 0.51 & 3.1 & 1 & 1,3,4 \\ TXS 1442+637 & 1.380 & $<1.3 $ & 23.90 & 27.51 & 0.05 & 2.0 & 1 & 1,3,4 \\ TXS 1030+611 & 1.401 & $<0.87$ & 23.56 & 27.53 & -0.37 & 5.3 & 2 & 1--4 \\ TXS 1010+350 & 1.410 & $<0.67$ & 23.94 & 27.41 & -0.43 & -- & 1 & 1 \\ TXS 2229+695 & 1.413 & $<1.7 $ & 24.47 & 27.15 & 0.64 & -- & 1 & 3,13 \\ TXS 0820+560 & 1.418 & $<0.34$ & 23.69 & 27.87 & -0.26 & -- & 2 & 1 \\ TXS 0805+410 & 1.418 & $<0.69$ & 23.32 & 27.40 & -0.07 & -- & 2 & 1,2 \\ TXS 0804+499 & 1.436 & $<0.95$ & 23.45 & 27.54 & 0.31 & -- & 2 & 1,2 \\ TXS 0145+386 & 1.442 & $<1.2 $ & 23.54 & 27.02 & 0.58 & 1.7 & 1 & 1,3,4 \\ TXS 0917+624 & 1.446 & $<0.79$ & 23.22 & 27.70 & 0.23 & 3.2 & 2 & 1--4 \\ JVAS J2311+4543 & 1.447 & $<2.2 $ & 22.82 & 26.90 & 0.76 & -- & 1 & 1,3 \\ S5 1058+72 & 1.460 & $<0.31$ & 24.27 & 27.89 & -0.21 & 1.9 & 1 & 1,3,4 \\ TXS 0859+470 & 1.470 & $<0.43$ & 23.54 & 28.17 & -0.23 & -- & 2 & 2 \\ TXS 2253+417 & 1.476 & $<0.71$ & -- & 27.82 & 0.30 & -- & 2 & -- \\ B3 1746+470 & 1.484 & $<4.5 $ & 23.43 & 27.00 & 0.54 & -- & 1 & 1,3 \\ TXS 0340+362 & 1.484 & $<3.0 $ & 22.75 & 27.19 & 0.35 & -- & 2 & 3 \\ TXS 1427+543 & 3.013 & $<0.44$ & 23.77 & 28.66 & -0.60 & -- & 1 & 2 \\ TXS 0800+618 & 3.033 & $<1.2 $ & 23.67 & 28.22 & -0.05 & -- & 2 & 1,3 \\ TXS 0642+449 & 3.396 & $<2.9 $ & 24.47 & 27.86 & 0.72 & 2.9 & 2 & 3,4,17 \\ TXS 0620+389 & 3.469 & $<0.20$ & 24.28 & 28.50 & -0.16 & 2.4 & 2 & 3,14,18 \\ TXS 0604+728 & 3.530 & $4.29 \pm 0.28$ & 23.96 & 28.64 & -0.38 & -- & 2 & 1,3 \\ TXS 0749+426 & 3.589 & $<0.76$ & 24.60 & 28.14 & 0.16 & 2.4 & 2 & 2,3,18 \\ \hline \hline \end{tabular} \end{center} \begin{tablenotes} \item[]~Notes to the table: \item[a]$^{a}$The inferred rest-frame 1216 \AA\ AGN luminosity, obtained by extrapolating from measurements in two nearby optical and/or ultraviolet bands. For two AGNs (indicated by a ``--" in this column), the UV luminosity could not be obtained as the AGN flux density is only available at a single optical waveband in the literature. \item[b]$^b$For sources with ``--" entries, the flux density is not known in the K-band; the (R$-$K) colour hence could not be obtained. \item[c]$^c$References for the associated \hii\ absorption searches : (1)~This paper; (2)~\citet{aditya16}; (3)~\citet[][]{gupta06}; (4)~\citet[][]{young1973}; (5)~\citet[][]{orienti06}; (6)~\citet[][]{vangorkom89}; (7)~\citet[][]{chandola13}; (8)~\citet[][]{dickey82}; (9)~\citet[][]{morganti09}; (10)~\citet[][]{peck99}; (11)~\citet[][]{vermeulen03}; (12)~\citet[][]{carilli98}; (13)~\citet[][]{curran13}; (14)~\citet[][]{gallimore99}; (15)~\citet[][]{kanekar08c}; (16)~\citet[][]{aditya17}. \item[d]$^d$References for the ultraviolet, optical, and infrared luminosity measurements, which were used to obtain the rest-frame 1216~\AA\ UV luminosity (following the procedure of \citealt[][]{curran10}), and the (R$-$K) colour: (1)~\citet[][]{bianchi14}; (2)~\citet[][]{abazajian09}; (2)~\citet[][]{monet03}; (4)~\citet[][]{cutri03}; (5)~\citet[][]{massaro04}; (6)~\citet[][]{cutri13}; (7)~\citet[][]{chen05b}; (8)~\citet[][]{howard04}; (9)~\citet[][]{odell78}; (10)~\citet[][]{raiteri09}; (11)~\citet[][]{urry00}; (12)~\citet[][]{zacharias04}; (13)~\citet[][]{veron10}; (14)~\citet[][]{souchay15}; (15)~\citet[][]{healey08}; (16)~\citet[][]{rao06}; (17)~\citet[][]{fedorov11}; (18)~\citet[][]{kuhn04}. \end{tablenotes} \end{table*} \subsection{Redshift evolution} \label{sec:redshift} Figure~\ref{fig:tau_vs_z} plots the velocity-integrated \hii\ optical depth, in logarithmic units, against AGN redshift, for the full sample of 92 sources. It is clear that our GMRT observations, especially at $1.1 \lesssim z \lesssim 1.5$, are sufficiently sensitive to detect \hii\ opacities lower than those of most of the detections of \hii\ absorption. Further, most \hii\ detections are concentrated at low redshifts, $z < 1$, with 13 detections at $z < 1$, and just 3 detections at $z > 1$ \citep[including the tentative detection at $z \approx 3.530$ towards TXS~0604+728;][]{aditya16}. The median redshift of the sample is $z_{\rm med} = 1.200$. Dividing the sample at this redshift into low-$z$ and high-$z$ sub-samples, the former has 13 detections and 33 non-detections of \hii\ absorption, whereas the latter has 3 detections (one of which is tentative) and 43 non-detections. The detection rates of \hii\ absorption (see Fig.~\ref{fig:det_frac}) are $28^{+10}_{-8}\%$ and $7^{+6}_{-4}\%$ \citep[again estimating the $1\sigma$ errors from Poisson statistics;][]{gehrels86} for the $z < z_{\rm med}$ and $z > z_{\rm med}$ sub-samples, respectively; the high-$z$ detection rate is even lower, $4^{+6}_{-3}\%$, when the tentative detection towards TXS~0604+728 is excluded from the sample. The high-$z$ sub-sample thus has a lower \hii\ detection rate, albeit only at $\approx 2.1\sigma$ significance. In addition to the \hii\ detection rates, we tested whether the distribution of the strength of the \hii\ absorption in AGN environments varies with redshift. For this purpose, we used survival analysis, in the {\sc asurv} package \citep[][]{isobe86}, to correctly include the upper limits on the \hii\ opacity. Within {\sc asurv}, the Peto-Prentice generalized Wilcoxon test finds that the null result that the velocity-integrated \hii\ optical depths of the low-$z$ and high-$z$ AGN sub-samples (again separated at $z_{\rm med}$) are drawn from the same distribution is ruled out at $3\sigma$ significance (increasing to $3.4\sigma$ significance when the tentative detection towards TXS~0604+728 is excluded from the sample). We thus find statistically significant evidence for redshift evolution in the strength of associated \hii\ absorption in a uniformly-selected AGN sample, with lower-redshift AGNs showing both a higher detection rate of \hii\ absorption and significantly higher integrated \hii\ optical depths. The above redshift dependence of the strength of associated \hii\ absorption could stem from a variety of reasons: (1)~less neutral hydrogen in high-$z$ AGN environments, implying lower \hi\ column densities, (2)~higher gas spin temperatures in high-$z$ environments, as has been seen in ``intervening'' galaxies towards AGNs, the damped Lyman-$\alpha$ absorbers \citep[e.g.][]{kanekar03,kanekar14}, or (3)~lower covering factors in the high-$z$ AGN sample, yielding a lower {\it observed} \hii\ optical depth. Unfortunately, in the case of associated \hii\ studies, we do not have direct estimates of the \hi\ column density, and hence cannot separate between the first two possibilities, a low \hi\ column density or a high gas spin temperature. We will initially consider the low covering factor hypothesis to account for the low detection rate of \hii\ absorption, before investigating whether AGN conditions might yield either a low gas content or a high spin temperature. \begin{figure} \includegraphics[scale=0.45]{fig2.pdf} \caption[]{The velocity-integrated \hii\ optical depth of the 92 CJF sources of the sample, plotted as a function of redshift. The 63 sources observed in the present GMRT survey are shown by squares, while the 29 sources from the literature are represented by triangles. Filled symbols indicate detections of \hii\ absorption, while open symbols indicate upper limits on the \hii\ optical depth. The dashed vertical line indicates the median redshift of the sample, $z_{\rm med} = 1.2$. \label{fig:tau_vs_z}} \end{figure} \begin{figure} \includegraphics[scale=0.45]{fig3.pdf} \caption[]{The detection rate of \hii\ absorption for the low-$z$ and high-$z$ sub-samples, separated at the median redshift. \label{fig:det_frac}} \end{figure} \subsection{Covering factor issues: The radio spectral index} \label{sec:tau_alpha} A possible cause for the lower strength of associated \hii\ absorption in high-$z$ AGNs is that the gas covering factor is systematically lower in the high-redshift systems. In such a scenario, the observed difference between the \hii\ absorption properties of the low-$z$ and high-$z$ sub-samples would arise not due to changes in the properties (e.g. \hi\ column density or spin temperature) of neutral hydrogen in AGN environments, but due to differences in the structure of the radio emission in the high-$z$ and low-$z$ AGNs of the sample. Specifically, if the low-frequency radio emission of high-$z$ AGNs arises primarily from extended structure (albeit still unresolved on GMRT baselines), which is not occulted by foreground gas clouds, the \hii\ optical depth estimated via the GMRT observations could be significantly lower than the true optical depth. For example, the radio emission of the low-$z$ AGNs at the redshifted \hii\ line frequency might predominantly arise from a compact radio core, while that of the high-$z$ AGNs might arise from either the radio jet or radio lobes. The simplest way of testing the above scenario is to measure the fraction of radio emission arising from the AGN core at, or close to, the redshifted \hii\ line frequency, via high-resolution VLBI imaging studies \citep[e.g.][]{kanekar09a,kanekar14}. This core fraction then gives a lower limit to the covering factor, under the assumption that the foreground gas clouds are likely to cover the core. Unfortunately, VLBI observations at frequencies $\lesssim 1$~GHz are technically challenging and are hence not available for most of the AGNs of our sample. An alternative approach to addressing the covering factor issue is based on the fact that the compact emission from the core tends to undergo synchrotron self-absorption and hence typically has an inverted or flat spectrum, while extended radio emission from the radio jet or the lobes tends to have a steep spectrum. If the radio emission of the high-$z$ AGNs at the redshifted \hii\ line frequency is dominated by the extended structure, yielding a low covering factor, one would expect these AGNs to have a systematically steeper spectral index at the \hii\ line frequency than the AGNs of the low-$z$ sub-sample. We note that, although our target AGNs have been uniformly chosen from the CJF sample, with flat spectral indices, $\alpha \geq -0.5$ \citep[][]{taylor96}, the CJF spectral index criterion is based on the AGN flux densities at two relatively high frequencies, 1.4~GHz and 4.85~GHz. It is hence possible that the radio emission at the low redshifted \hii\ line frequency is dominated by steep-spectrum extended structure, rather than by the flat- or inverted-spectrum radio core. We will hence use the AGN's radio spectral index $\alpha_{\rm 21-cm}$ at the redshifted \hii\ frequency as a proxy for the compactness of the AGN. A flat or an inverted spectrum near the redshifted \hii\ line frequency ($\alpha_{\rm 21-cm} \gtrsim 0$) would indicate a core-dominated source, and a relatively high covering factor ($f \approx 1$), whereas a steep spectrum ($\alpha_{\rm 21-cm} \lesssim -0.7$) would indicate extended radio structure and a possibly low covering factor ($f \ll 1$). If the measured \hii\ optical depths are found to depend on the spectral index, or if the high-$z$ AGN sub-sample has a systematically steeper spectral index than the low-$z$ sample, it would suggest that a low covering factor may be the cause of the lower observed \hii\ optical depths at high redshifts. \begin{figure*} \includegraphics[scale=0.4]{fig4a.pdf} \includegraphics[scale=0.4]{fig4b.pdf} \caption{[A]~Left panel: The integrated \hii\ optical depth of the 92 CJF sources plotted against the spectral index, $\alpha_{\rm 21-cm}$, around the redshifted \hii\ line frequency; the dashed vertical line indicates the median spectral index, $\alpha_{\rm 21-cm,med} = -0.015$. [B]~Right panel: The spectral index $\alpha_{\rm 21-cm}$, plotted against AGN redshift; the dashed vertical line indicates the median redshift, $z_{\rm med} = 1.2$, while the dashed horizontal line indicates the median spectral index, $\alpha_{\rm 21-cm} = -0.015$. In both panels, the 63 AGNs of our survey are shown as squares, while the 29 literature sources are shown as triangles. Filled symbols represent detections, while open symbols represent upper limits on the integrated \hii\ optical depth. See main text for discussion. \label{fig:alpha}} \end{figure*} Figure~\ref{fig:alpha}[A] shows the integrated \hii\ optical depths of the 92 CJF sources of our sample plotted against $\alpha_{\rm 21-cm}$. While there are clearly a few sources with $\alpha_{\rm 21-cm} \lesssim -0.5$, most of the AGNs are seen to have flat radio spectra, $\alpha_{\rm 21-cm} \approx 0$. Indeed, the median spectral index of the sample, $\alpha_{\rm 21-cm,med} = -0.015$, is very close to zero, indicating that the sample is dominated by compact objects, with $\alpha_{\rm 21-cm} > -0.5$. A Peto-Prentice two-sample test finds that the distributions of the \hii\ optical depths of the two sub-samples, separated at the median $\alpha_{\rm 21-cm}$, are consistent (within $\approx 1.3\sigma$ significance) with the null hypothesis of being drawn from the same underlying distribution. We thus find no evidence for a dependence of the strength of the associated \hii\ absorption on the AGN spectral index. Figure~\ref{fig:alpha}[B] shows the low-frequency AGN spectral index $\alpha_{\rm 21-cm}$ plotted against redshift; no trend is apparent in the figure. A Gehan-Wilcoxon two-sample test comparing the distributions of $\alpha_{\rm 21-cm}$ values of the low-$z$ and the high-$z$ sub-samples, separated at $z_{\rm med} = 1.2$, finds that the data are consistent (within $\approx 1.2\sigma$ significance) with the null hypothesis of being drawn from the same underlying distribution. We thus find no evidence for a systematic difference between the spectral indices of the AGNs of the high-$z$ and the low-$z$ sub-samples. In summary, we find no statistically significant evidence either that the strength of the \hii\ absorption depends on the low-frequency AGN spectral index or that the low-frequency spectral index varies systematically with redshift. Both of these would have been expected if low covering factors are the cause of the weaker \hii\ absorption observed in the high-$z$ AGN sub-sample. We hence conclude that it is unlikely that the observed differences in the \hii\ absorption properties of the low-$z$ and high-$z$ AGN sub-samples can be explained by covering factor issues. \subsection{The AGN colour: Evidence for dust reddening?} \label{sec:dust} \begin{figure} \includegraphics[scale=0.4]{fig5.pdf} \caption[]{The integrated \hii\ optical depth, in logarithmic units, plotted against the (R$-$K) colour for the 58 AGNs of the sample with NIR photometry. The dashed vertical line indicates the median colour, $\rm (R$-$K) = 2.38$. The sources from our survey and the literature are shown as squares and triangles, respectively. Detections of \hii\ absorption are shown as filled symbols and non-detections as open symbols. See main text for discussion. \label{fig:tau_rk}} \end{figure} For red quasars, the high extinction at optical wavebands is believed to be caused by dust extinction \citep[e.g.][]{webster95}. However, it has also been suggested in the literature that not all red quasars are dusty systems \citep[e.g.][]{benn98}. Earlier \hii\ studies have yielded ambiguous results: for example, \citet[][]{carilli98} detected strong \hii\ absorption in four out of five red quasars at intermediate redshifts ($z \approx 0.7$), suggesting a high \hi\ column consistent with the dust obscuration hypothesis. Strong molecular absorption has also only been seen in very red AGNs \citep[e.g.][]{wiklind94,wiklind96,kanekar02,kanekar05}. However, later \hii\ absorption studies have found no significant correlation between the detectability of associated \hii\ absorption and the AGN colour \citep[e.g.][]{curran08,aditya16}, suggesting that the reddening may have other causes besides dust obscuration. Our detection of associated \hii\ absorption towards TXS~1456+375 is consistent with the dust-reddening hypothesis. It has long been known that the Galactic extinction at 5500~\AA\ is proportional to the total hydrogen column density, with $A_{v}=6.29 \times 10^{-22} \left[ N_{\rm HI} + 2N_{\rm H_2} \right]$ \citep[e.g.][]{savage77}. This is understood as arising due to reddening produced by the large amounts of dust associated with a high hydrogen column. \citet{webster95} hence suggested that ``red'' quasars, with a steeper spectral index between the optical and NIR wavebands than typical quasars, are likely to acquire their redder colours due to dust extinction. Consistent with this hypothesis, \citet{carilli98} found that 80\% of red quasars showed either associated or intervening \hii\ absorption, while only 11\% of optically-selected Mg{\sc ii} absorbers showed radio absorption. Similarly, all five of the known redshifted radio molecular absorbers at $z > 0.2$ have background quasars with extremely red colours, $\rm (R$-$K) > 4$ \citep{wiklind94,wiklind95,wiklind96,wiklind97,kanekar02,kanekar05,curran06b,curran08}. Our detection of associated \hii\ absorption towards TXS~1456+375 is consistent with the dust-reddening hypothesis. However, it has also been noted in the literature that not all red quasars appear to be dusty systems \citep[e.g.][]{benn98}. Indeed, recent searches for associated \hii\ absorption in red AGNs have not been very successful \citep[e.g.][]{yan16}, while no significant correlation has been found between the detectability of associated \hii\ absorption and the AGN colour \citep[e.g.][]{curran08,aditya16}, suggesting that the reddening may have other causes besides dust obscuration. We note that the sample of \citet{carilli98} was very small (five systems), and limited to low redshifts, $z \lesssim 0.7$. In this section, we test the hypothesis that associated \hii\ absorption is more likely to arise in red AGNs, by examining our sample for a correlation between the strength of the \hii\ absorption with AGN (R$-$K) colour. Unfortunately, NIR photometry was available for 58 AGNs of the sample and so the present analysis is restricted to these 58 systems. At the outset, we note that the choice of a flat-spectrum sample is likely to be biased against the reddest sightlines. In standard unification schemes \citep[e.g.][]{urry95}, a flat AGN spectrum is expected to mostly arise from sightlines closer to orthogonal to the obscuring torus. Such sightlines are unlikely to be affected by extinction from dust in the torus, and so we do not expect very large (R$-$K) values, $\rm (R$-$K) \gtrsim 5$, in our sample. Fig.~\ref{fig:tau_rk} shows the integrated \hii\ optical depth plotted versus the (R$-$K) colour for the 58 AGNs with NIR photometry. Of the 12 AGNs with detections of \hii\ absorption, five have relatively red colours, (R$-$K)~$> 3$. Further, three of these five systems lie at the top right of the figure, indicating both red colours and high integrated \hii\ optical depths; indeed, these three AGNs have $(R$-$K) > 5$, comparable to the colours of the AGNs that show molecular absorption! At least for these systems, the red colours are likely to arise due to the presence of high columns of gas and associated dust at the AGN redshift. However, we also note that there are four AGNs that show \hii\ absorption at relatively low (R$-$K) values, (R$-$K)~$\lesssim 1.3$, and that one of these has the second-highest integrated \hii\ optical depth of the full sample. To test the dependence of the strength of the \hii\ absorption on the (R$-$K) colour, we divided the sample of 58 systems at the median (R$-$K) value, $\rm (R$-$K)= 2.38$, and carried out two-sample tests on the high-(R$-$K) and low-(R$-$K) sub-samples. A Peto-Prentice two-sample test for censored data finds that the null hypothesis that the sub-samples are drawn from the same underlying distribution is rejected at only $1.4\sigma$ significance. The two sub-samples are thus consistent with being drawn from the same distribution, and we find no statistically significant evidence for a dependence of the strength of \hii\ absorption on the (R$-$K) colour of the AGN. Thus, while three of the four highest integrated \hii\ optical depths do arise from the reddest AGNs, our results from the CJF sample do not provide support for the dust-reddening hypothesis. \subsection{Effects of the AGN luminosity} \label{sec:uv_radio_lum} \begin{figure*} \includegraphics[scale=0.4]{fig6a.pdf} \includegraphics[scale=0.4]{fig6b.pdf} \caption{The AGN luminosities of the CJF sources of the sample in [A]~(left panel) at rest-frame UV 1216~\AA\ (87 sources) and [B]~(right panel) rest-frame radio~1.4~GHz (92 sources), plotted, in logarithmic units, against the AGN redshift. The dashed vertical lines indicate the median redshift, $z_{\rm med}$, while the dashed horizontal lines indicate the median UV 1216~\AA\ and radio 1.4~GHz luminosities ($L_{UV,med} = 10^{22.64}$~W~Hz$^{-1}$ and $L_{1.4\;GHz,med} = 10^{27.17}$~W~Hz$^{-1}$, respectively). The squares and triangles represent, respectively, the sources from our survey and the literature. Filled and open symbols represent, respectively, \hii\ detections and upper limits on the integrated \hii\ optical depth. See text for discussion.} \label{fig:lum_z} \end{figure*} As mentioned briefly in Section~\ref{sec:sample}, a high AGN luminosity can adversely affect the strength of the associated \hii\ absorption \citep[e.g.][]{curran08,curran13}. This is because a high UV luminosity can cause ionization of the nearby neutral hydrogen, thus reducing the \hi\ column density, and can also affect the spin temperature \citep[which is coupled to the colour temperature of the Lyman-$\alpha$ radiation field; ][]{wouthuysen52,field58,field59}. Similarly, a high AGN rest-frame 1.4~GHz luminosity can raise the gas spin temperature \citep[][]{field58,field59}. Both a decrease in the \hi\ column density and a raising of the spin temperature would reduce the strength of the \hii\ absorption. Thus, if the higher-redshift AGNs of our sample tend to have higher UV and/or radio luminosities, this would provide an alternative explanation of the weaker \hii\ absorption in the high-$z$ sub-sample. Figs.~\ref{fig:lum_z}[A] and [B] show, respectively, the rest-frame UV 1216~\AA\ luminosity and the rest-frame 1.4~GHz radio luminosity of the sources of the CJF sample plotted against redshift. In both panels, the vertical dashed line is at the median redshift of the sample, while the horizontal dashed line is at the median luminosity. It is clear that the high-$z$ AGNs of the sample typically lie close to or above the median luminosity (in both UV and radio wavebands), while the low-$z$ AGNs mostly lie below the median luminosity. Again dividing the sample at $z_{\rm med} = 1.2$, a Gehan-Wilcoxon test finds that the null hypothesis that the rest-frame UV and radio luminosities of the two sub-samples are drawn from the same distribution is rejected, respectively, at $\approx 7.7\sigma$ and $\approx 7.8\sigma$ significance. As such, it is clear that the sample contains a strong bias towards higher UV and radio luminosities at high redshifts. We note, in passing, that the rest-frame UV $1216$~\AA\ and radio 1.4~GHz luminosities of the sources in our sample are strongly correlated (at $\approx9\sigma$ significance, via a Kendall-tau test). \begin{figure*} \includegraphics[scale=0.4]{fig7a.pdf} \includegraphics[scale=0.4]{fig7b.pdf} \caption[]{The integrated \hii] optical depths of the AGNs of the our sample plotted against [A]~(left panel) the rest-frame UV 1216~\AA\ luminosity and [B]~(right panel) the rest-frame radio 1.4~GHz luminosity, with all quantities in logarithmic units. The dashed vertical lines indicate the median UV 1216~\AA\ and radio 1.4~GHz luminosities ($L_{UV,med} = 10^{22.64}$~W~Hz$^{-1}$ and $L_{1.4\;GHz,med} = 10^{27.17}$~W~Hz$^{-1}$, respectively). The sources from our survey and the literature are represented by squares and triangles, respectively. Filled and open symbols represent, respectively, detections and upper limits on the \hii\ optical depth. \label{fig:tau_lum}} \end{figure*} It thus appears that the apparent redshift evolution in the strength of the associated \hii\ absorption might also arise due to differences in the UV and/or radio luminosities of the AGNs of the low-$z$ and the high-$z$ sub-samples. Figs.~\ref{fig:tau_lum}[A] and [B] plot the integrated \hii\ optical depth versus, respectively, the rest-frame UV 1216~\AA\ luminosity and the rest-frame radio 1.4~GHz luminosity (with all quantities in logarithmic units). The dashed vertical lines in the two figures indicate the median luminosities, $L_{UV,med} = 10^{22.64}$~W~Hz$^{-1}$ and $L_{1.4\;GHz,med} = 10^{27.17}$~W~Hz$^{-1} $, respectively. It is clear that most of the detections of \hii\ absorption lie in the low-luminosity halves of the two figures; further, the integrated \hii\ optical depths towards low-luminosity AGNs appear higher than the typical $3\sigma$ upper limits on the integrated \hii\ optical depths towards AGNs with high luminosities. Formally, a Peto-Prentice two-sample test (for censored data) finds that the null hypotheses that the \hii\ optical depth distributions of the low-luminosity and high-luminosity sub-samples are drawn from the same distribution is rejected at, respectively, $\approx 3.6\sigma$ and $\approx 3.3\sigma$ significance, for the rest-frame UV $1216$~\AA\ and rest-frame radio 1.4~GHz luminosities. We thus find statistically significant evidence for a dependence of the strength of associated \hii\ absorption in the CJF sample on both redshift and AGN luminosity in the rest-frame UV 1216~\AA\ and radio 1.4~GHz wavebands, but not on the low-frequency radio spectral index or the (R$-$K) colour. Weaker \hii\ absorption is obtained at higher redshifts and higher radio and UV luminosities. Unfortunately, there is a strong correlation between AGN luminosity and redshift in the AGNs of our sample (see Figs.~\ref{fig:lum_z}[A] and [B]), due to which it is not currently possible to break the degeneracy between redshift and luminosity, and identify the primary cause, if any, for the differences in the strength of the \hii\ absorption. Searches for \hii\ absorption in either a low-luminosity AGN sample at high redshifts, or a high-luminosity sample at low redshifts would be required to break the present degeneracy. We have used the GMRT to carry out a search for associated \hii\ absorption in 50 flat-spectrum AGNs, selected from the CJF sample. The data on ten AGNs were rendered unusable by RFI. We obtained new detections of \hii\ absorption in two sources, at $z = 0.229 $ towards TXS~0003+380 and $z =0.333$ towards TXS~1456+375, and also confirmed an earlier detection \citep[by][]{curran13} of \hii\ absorption at $z = 1.277$ towards TXS~1543+480. The measured velocity-integrated \hii\ optical depths towards the above three AGNs lie in the range $\int \tau d{\rm V} \approx 1.9 - 9.6$~\kms, implying \hi\ column densities of $\approx (3.5 - 17.5) \times 10^{20}$~cm$^{-2}$, for an assumed spin temperature of $100$~K and covering factor of unity. For the remaining 37 AGNs, the $3\sigma$ upper limits on the integrated \hii\ optical depth range from $0.3 - 14$~\kms, with a median value of $\approx 0.97$~\kms. Our full sample of CJF sources with searches for redshifted \hii\ absorption consists of 92 AGNs, 63 from our survey \citep[including 24 sources from][]{aditya16,aditya17} and 29 from the literature. This is currently the largest sample of uniformly-selected AGNs with searches for associated \hii\ absorption, with 16 \hii\ detections and 76 non-detections, at redshifts $0.01-3.6$, and with a median redshift of $z_{\rm med} = 1.2$. We find that both the strength and the detectability of \hii\ absorption appear higher at low redshifts, $z < z_{\rm med}$. The detection rate of \hii\ absorption is $28^{+10}_{-8}$\% for the low-$z$ AGN sub-sample (with $z < 1.2$), but only $7^{+6}_{-4}$\% for the high-$z$ sub-sample (with $z > 1.2$). While the difference in detection rates has only $\approx 2.1\sigma$ significance, a Peto-Prentice two-sample test on the velocity integrated \hii\ optical depths finds that the null hypothesis that the low-$z$ and high-$z$ sub-samples are drawn from the same distribution is ruled out at $\approx 3\sigma$ significance. We thus obtain statistically-significant evidence for redshift evolution in the strength of associated \hii\ absorption in the Caltech-Jodrell Flat-spectrum AGN sample. However, we also found evidence for a significant bias in the intrinsic luminosities of the AGNs of our sample, with the high-$z$ AGNs having higher rest-frame 1216~\AA\ UV and 1.4~GHz radio luminosities. Examining the dependence of the strength of the \hii\ absorption on AGN luminosity, the null hypothesis that the velocity-integrated \hii\ optical depths of the high-luminosity and low-luminosity AGNs arise from the same distribution is ruled out at $\approx (3.3-3.6)\sigma$ significance in a Peto-Prentice two-sample test (for the 1216~\AA\ UV and 1.4~GHz radio luminosities). We also examined the possibility that the lower strength of \hii\ absorption in high-$z$ AGNs might arise due to a typically lower covering factor for the high-$z$ sub-sample. This could occur if the radio continuum at the redshifted \hii\ line frequency for the high-$z$ sub-sample is dominated by extended emission. We used the AGN spectral index around the redshifted \hii\ line frequency as a proxy for source compactness, since extended emission is expected to have a steep spectrum. We find no evidence in two-sample tests that the strength of the \hii\ absorption depends on AGN spectral index, or that the spectral index itself depends on the AGN redshift. It is hence unlikely that the observed weakness of the \hii\ absorption in high-$z$ AGNs arises due to a low covering factor. We find no statistically-significant evidence that the strength of \hii\ absorption depends on AGN colour. However, five of the 12 AGNs with \hii\ detections and estimates of the (R$-$K) colour show relatively red colours, with (R$-$K)~$>3$. Three of these systems have both high integrated \hii\ optical depths and (R$-$K)~$> 5$; for these, the red colour is likely to arise due to dust at the AGN redshift. The above results are consistent with those from our pilot GMRT \hii\ absorption survey of the CJF sample \citet{aditya16}, but are now based on a significantly larger AGN sample (92 AGNs against 52 in \citealt{aditya16}), and with more than half the sample at $z > 1$. The median redshift of the present sample ($z_{\rm med} = 1.2$) is also significantly higher than that ($z_{\rm med} = 0.76$) of \citet{aditya16}. In summary, the strength of associated \hii\ absorption in the CJF AGN sample appears to depend on both redshift and AGN luminosity, with weaker \hii\ absorption at high redshifts and high luminosities. This could arise due to (1)~lower amounts of neutral gas in high-$z$, high-luminosity AGN environments, due to either redshift evolution or ionization of the \hi\ by the far-UV AGN radiation, or (2)~higher spin temperatures in high-$z$, high-luminosity AGN environments, due to either a preponderance of warm neutral gas around high-$z$ AGNs or spin temperatures greater than the kinetic temperature due to the high AGN UV/radio luminosity. Unfortunately, the luminosity bias in our sample, with the higher-luminosity AGNs located at higher redshifts, implies that the present dataset does not allow us to distinguish between the above possibilities to identify whether redshift or AGN luminosity is the primary driving factor in determining the strength of the associated \hii\ absorption; this will be the focus of future studies. | 18 | 8 | 1808.03280 |
1808 | 1808.03763_arXiv.txt | Radioactive energies from unstable nuclei made in the ejecta of neutron star mergers play principal roles in powering kilonovae. In previous studies power-law-type heating rates (e.g., $\propto t^{-1.3}$) have frequently been used, which may be inadequate if the ejecta are dominated by nuclei other than the $A\sim 130$ region. We consider, therefore, two reference abundance distributions that match the $r$-process residuals to the solar abundances for $A \ge 69$ (light trans-iron plus $r$-process elements) and $A \ge 90$ ($r$-process elements). Nucleosynthetic abundances are obtained by using free-expansion models with three parameters: expansion velocity, entropy, and electron fraction. Radioactive energies are calculated as an ensemble of weighted free-expansion models that reproduce the reference abundance patterns. The results are compared with the bolometric luminosity ($>$ a few days since merger) of the kilonova associated with GW170817. We find that the former case (fitted for $A \ge 69$) with an ejecta mass $0.06\, M_\odot$ reproduces the light curve remarkably well including its steepening at $\gtrsim 7$~days, in which the mass of $r$-process elements is $\approx 0.01\, M_\odot$. Two $\beta$-decay chains are identified: $^{66}$Ni$\rightarrow^{66}$Cu$\rightarrow^{66}$Zn and $^{72}$Zn$\rightarrow^{72}$Ga$\rightarrow^{72}$Ge with similar halflives of parent isotopes ($\approx 2$~days), which leads to an exponential-like evolution of heating rates during 1--15~days. The light curve at late times ($> 40$~days) is consistent with additional contributions from the spontaneous fission of $^{254}$Cf and a few Fm isotopes. If this is the case, the event GW170817 is best explained by the production of both light trans-iron and $r$-process elements that originate from dynamical ejecta and subsequent disk outflows from the neutron star merger. | \label{sec:introduction} The discovery of a binary neutron star (NS) merger as the source of gravitational wave signals confirmed by advanced LIGO/Virgo \citep[on August 17, 2017; GW170817,][]{Abbott+2017}, followed by the detection of an electromagnetic counterpart across the entire wavelength range, has opened a window of the ``multi-messenger astronomy" \citep[see, e.g.,][]{Metzger2017}. Among the latter, the observations of electromagnetic emission in optical and near-infrared ranges \citep[kilonova,][]{Li1998, Metzger2010} provided us with numerous clues to understanding the origin of heavy elements such as gold and uranium made by the rapid neutron-capture process \citep[$r$-process, see][for a recent review]{Thielemann2017}. The history of the $r$-process study goes back to the seminal works by \citet{Burbidge1957} and \citet{Cameron1957}. \citet{Burbidge1957} supposed a sort of stellar explosions including core-collapse supernovae and Type~Ia supernovae be the astrophysical site of the $r$-process. Prior to these works, in \citet{Burbidge1956} a concept of the $r$-process already has been presented by associating the light curves of Type~Ia supernovae decaying in an exponential manner with the spontaneous fission of a transuranic species $^{254}$Cf (with a halflife of 60.5 days). It was before the identification of the decay chain $^{56}$Ni$\rightarrow^{56}$Co$\rightarrow^{56}$Fe \citep{Colgate1969, Arnett1979} that actually powers Type~Ia supernovae. Since then, a main focus of the study of $r$-process sites has been placed on core-collapse supernovae including prompt explosions \citep{Schramm1973, Sato1974, Hillebrandt1976, Sumiyoshi2001, Wanajo2003} and more recently neutrino-driven explosions \citep{Meyer1992, Woosley1994, Witti1994, Qian1996, Cardall1997, Otsuki2000, Wanajo2001, Thompson2001}. However, recent studies exclude the former explosion mechanism \citep[e.g.,][]{Kitaura2006, Janka2012} and suggest the latter be sources of only light trans-iron elements \citep[e.g.,][]{Wanajo2011, Wanajo2013}. Currently, only a magneto-rotationally induced explosion mechanism remains a viable possibility among the scenarios of core-collapse supernovae \citep{Winteler2012, Nishimura2015, Nishimura2017, Moesta2017, Halevi2018}. An alternative scenario, the decomposition of neutron-rich material from merging binary neutron stars (NSs, or a NS and a black hole) also was proposed as the sources of $r$-process elements \citep{Lattimer1974, Symbalisty1982, Eichler1989, Meyer1989}. Early studies of the dynamical ejecta (resulting from tidal torque and shock heating) of NS mergers pointed the robust production of heavy $r$-process nuclei ($A > 120$) by a fission recycling in extremely neutron-rich environments \citep{Freiburghaus1999, Goriely2011, Korobkin2012, Bauswein2013}. More recent studies indicate, however, the production of all the $r$-process species ($A \ge 90$) in the ejecta with a wide range of neutron-richness owing to weak interactions \citep[electron/positron and neutrino captures on free nucleons,][]{Wanajo2014, Sekiguchi2015, Sekiguchi2016, Goriely2015, Radice2016, Papenfort2018}. The post-merger outflows from the accretion disk orbiting around a remnant (a massive NS or a black hole) also were suggested as a site of the $r$-process \citep{Ruffert1999, Metzger2008, Surman2008, Wanajo2012}. Recent studies indicate various possibilities such that the disk outflows are modestly \citep[e.g.,][]{Just2015, Fujibayashi2018} or very \citep[e.g.,][]{Wu2016, Siegel2017, Fernandez2018} neutron-rich. Latest studies of Galactic chemical evolution imply that $r$-process-enhanced halo stars reflect the nucleosynthetic yields from single NS merger events \citep{Ishimaru2015, Beniamini2016, Safarzadeh2017, Ojima2018}. Spectroscopic analyses of Galactic halo stars also have been providing us with several important clues. The remarkable agreement of the abundance distributions of $r$-process-enhanced halo stars with those of the solar $r$-process component indicates a robust and single origin of $r$-process elements (e.g., CS~22092-052, \citealt{Sneden2003}; CS~31082-001, \citealt{Siqueira2013}). The stars observed in a recently discovered ultra-faint dwarf galaxy, Reticulum~II, also exhibit solar-like $r$-process abundance patterns \citep{Ji2016, Ji2018}. However, such a remarkable agreement appears not to extend to the low-$Z$ ($< 50$) and high-$Z$ ($> 80$) ends; abundance variations of a factor of several can be seen in lighter elements \citep[e.g., Sr, Y, and Zr,][]{Sneden2008} as well as in actinides \citep[e.g., Th,][]{Holmbeck2018} with respect to those in between (e.g., Eu). More seriously, there are few measurable lines of light trans-iron elements ($30 < Z < 40$) that compose the low-mass side of the $r$-process residuals to the solar system abundances \citep[e.g.,][]{Goriely1999}. Currently, only Ga and Ge have been measured from the space for only a few halo stars \citep{Sneden2008}. The identification of a kilonova (AT~2017gfo or SSS17a) associated with GW170817 has revealed the production of elements beyond iron in the NS merger \citep[e.g.,][]{Arcavi2017, Chornock2017, Cowperthwaite2017, Kasen2017, Kasliwal2017, Nicholl2017, Pian2017, Smartt2017, Tanaka2017, Tanvir2017}. While the early ``blue" emission indicates a lanthanide-free ($Z < 57$) component in the ejecta \citep{Metzger2014, Kasen2015, Tanaka2018}, the late-time ($>$ a few days) emission in red-optical and near-infrared wavelengths confirms the presence of freshly synthesized lanthanides that have high opacities \citep{Banes2013, Kasen2013, Tanaka2013}. However, the inferred mass fraction of lanthanides and heavier in the ejecta is only $\approx 10^{-4}$--$10^{-2}$ \citep[e.g.,][]{Arcavi2017, Chornock2017, Nicholl2017, Waxman2017}. It is questionable, therefore, if the merger made heavy $r$-process elements such as gold and uranium. Moreover, such photometric analyses alone cannot discriminate between lanthanides and heavier elements. Another problem is the large amount of ejecta from the merger; the inferred masses of the blue and red components are, respectively, $\approx 0.01$--$0.02\, M_\odot$ with the outflow velocity of $\approx 0.2\, c$ ($c$ is the speed of light) and $\approx 0.04\, M_\odot$ with $\approx 0.1\, c$ \citep[e.g.,][]{Cowperthwaite2017, Nicholl2017}. The total mass $\approx 0.05$--$0.06\, M_\odot$ is too large to be fulfilled by the dynamical ejecta of $\lesssim 0.01\, M_\odot$ \citep[][]{Hotokezaka2013, Bauswein2013, Sekiguchi2016, Radice2016}. The disk outflows may eject more material of $\sim 0.01$--$0.1\, M_\odot$ but with smaller velocity \citep[$\sim 0.05\, c$,][]{Dessart2009, Metzger2014, Just2015, Siegel2017, Shibata2017, Fujibayashi2018}. Note that most of the above estimates for the kilonova ejecta were based on the power-law-type heating rates \citep[e.g., $\approx 2\times 10^{10}\, t^{-1.3}$~erg~g$^{-1}$~s$^{-1}$, where $t$ is time in days,][]{Metzger2010, Wanajo2014} originating from the decaying radioactivities with $A\sim 130$. In this paper we revisit the issue of the radioactive heating rates in NS merger ejecta, which is supposed to be the first of the series of papers that explore the physical conditions for the $r$-process by using a multi-component free-expansion model described in \S~\ref{sec:model}. Nucleosynthetic abundances are obtained by using free-expansion models that cover a wide range of parameters (expansion velocity, entropy, and electron fraction). The heating rates are then calculated as an ensemble of free-expansion models with their weighted abundances, which fit the $r$-process residuals to the solar system abundances \citep{Goriely1999} for two cases: a) $A\ge 69$ and b) $A\ge 90$ (\S~\ref{sec:reference}). Obviously, the choice of reference abundance distributions are not unique; these two cases are taken for simplicity, in which the nuclei of $A \sim 70$ and 130, respectively, play dominant roles for radioactive heating. The resultant heating rates are presented in \S~4 with discussion on the contributions from $\beta$-decay, $\alpha$-decay, and fission. In \S~5, the heating rates for the two cases by adopting the thermalization efficiencies in \citet{Barnes2016} are compared with the kilonova light curve of the NS merger GW170817. Summary and conclusions follow in \S~6. | \label{sec:summary} Multi-component free-expansion model (mFE) was constructed to investigate various aspects of physical conditions relevant for the $r$-process. This paper was supposed to be the first of a series of papers with an emphasis on radioactive heating rates that power a kilonova, an electromagnetic counterpart of the gravitational signals from a NS merger. Each free expansion model (FE) consists of a three-parameter suite: expansion velocity $v/c\, (= 0.05$--0.30), entropy $S\, (= 10$--35 in units of $k_\mathrm{B}/\mathrm{nuc}$), and electron fraction $Y_\mathrm{e}\, (= 0.01$--0.50). A mFE was defined as an ensemble of FEs that fitted a given reference abundance distribution. The $r$-process residuals (by subtracting the $s$-process component) to the solar system abundances \citep{Goriely1999} was employed as the reference, in light of the robust (solar-$r$-process like) abundance patterns in $r$-process-enhanced stars \citep[at least in the range $50 < Z < 80$, e.g.,][]{Sneden2008}. Two models were considered: a) mFE-a with a reference of the full range of the residuals ($A \ge 69$) with both light trans-iron and $r$-process elements and b) mFE-b with $r$-process elements only ($A \ge 90$). For mFE-b, the fitting of FEs to the reference resulted in wide ranges of $S$ and $Y_\mathrm{e}$ as found in numerical simulations of dynamical ejecta (\citealt{Wanajo2014, Sekiguchi2015, Sekiguchi2016}; see also a similar trend in the magneto-hydrodynamic simulations of accretion disks by \citealt{Siegel2017, Fernandez2018}). In contrast, mFE-a composed of narrow ranges of $S\approx 10$ and $Y_\mathrm{e} \approx 0.4$ that were in good agreement with a recent simulation of disk outflows (\citealt{Fujibayashi2018}; see also \citealt{Lippuner2017}). As such physical conditions led to NSE, the nucleosynhthetic yields extended down to $^{48}$Ca, a nuclide whose astrophysical origin was unknown \citep{Hartmann1985, Meyer1996, Woosley1997, Wanajo2013b}. While the obtained heating rate for mFE-b exhibited a power-law-type temporal evolution owing to $\beta$-decay (of $A\sim 130$) as found in previous works \citep[$\approx 2 \times 10^{10}\, t^{-1.3}$~erg~g$^{-1}$~s$^{-1}$, e.g.,][]{Metzger2010, Wanajo2014}, that for mFE-a indicated rather an exponential-type evolution during $t\approx 1$--15~days. Two $\beta$-decay chains relevant for the heating were identified: $^{66}$Ni$\rightarrow^{66}$Cu$\rightarrow^{66}$Zn and $^{72}$Zn$\rightarrow^{72}$Ga$\rightarrow^{72}$Ge with both having similar halflives ($\approx 2$~days) of the parent isotopes made in NSE. Contributions from fission and $\alpha$-decay became non-negligible at late times ($> 10$~days) as pointed out by \citet{Hotokezaka2016, Barnes2016, Zhu2018}. Obtained heating rates with given ejecta mass $M_\mathrm{ej}$ and ejecta velocity $v_\mathrm{ej}$ \citep[both affect the thermalization efficiency,][]{Barnes2016} were compared with the bolometric light curve of the kilonova (AT~2017gfo or SSS17a) associated with GW170817. It was found that the light trans-iron dominant model mFE-a with $M_\mathrm{ej}/M_\odot = 0.06$ and $v_\mathrm{ej}/c = 0.1$ reproduced the bolometric light curve remarkably well. A steepening of the light curve at $\gtrsim 7$~days \citep{Waxman2017} was indicative of the dominance of light trans-iron nuclei including $^{66}$Ni and $^{72}$Zn, rather than $r$-process products. Late-time estimates at several 10 days \citep{Villar2018} could also be explained with the $\beta$-decay heating and probably those from fission and $\alpha$-decay. The pure $r$-process model (mFE-b with $M_\mathrm{ej}/M_\odot = 0.04$) did not account for the light curve at late times because of its robust power-law like decay of the heating rate. Note that our models account for the gross nature of the kilonova light curve but its early blue component that may originate from a high-latitude dynamical ejecta \citep[e.g.,][]{Shibata2017} or the wind ejecta from a strongly magnetized hypermassive NS \citep{Metzger2018}. In conclusion, the ejecta from the NS merger GW170817 was dominated ($\approx 0.05\, M_\odot$) by light trans-iron elements ($A < 90$) with a fraction ($\approx 0.01\, M_\odot$) of $r$-process elements ($A \ge 90$). Along with the adopted velocity ($v_\mathrm{ej}/c = 0.1$), our conclusion is consistent with an interpretation that the $r$-process elements come from the dynamical ejecta of a NS merger and light trans-iron elements from the subsequent disk outflows \citep[e.g.,][]{Shibata2017}. Magneto-hydrodynamic mass ejection from an accretion disk may also be a viable mechanism for this event \citep{Siegel2017, Fernandez2018}. Although the inferred $M_\mathrm{ej}$ and $v_\mathrm{ej}$ are similar to the literature values, it is emphasized that the principal radioactive energy sources are light trans-iron elements, not $r$-process elements as suggested in previous works. A word of caution is needed; this study itself cannot answer satisfactorily to a question: ``how much $r$-process elements were made?". This is due to the fact that the choice of a reference abundance distribution was arbitrary and our conclusion strongly relied upon the abundances of only two radioactive isotopes $^{66}$Ni and $^{72}$Zn. The inferred mass of $r$-process elements ($\approx 0.01$; Table~2) merely reflects our choice of the reference abundance distribution, i.e., the $r$-process residuals to the solar system abundances. Comparison of our model (mFE-a) with the observed bolometric luminosity only ensures the production of light trans-iron isotopes $^{66}$Ni and $^{72}$Zn with the amounts presented in Table~\ref{tab:properties} (5th and 6th columns)\footnote{The light curve cannot disentangle the contributions of two decay chains from $^{66}$Ni and $^{72}$Zn because of their similar lifetimes. Measurements of $\gamma$-ray lines (97.8~keV and 145~keV/2.77~MeV, respectively) from a future nearby event may enable us to directly determine their abundances \citep[see, e.g.,][]{Hotokezaka2016}.} but the abundance distribution in Figure~\ref{fig:nuclei} (top-left)\footnote{For instance, an addition of ejecta with $Y_\mathrm{e} = 0.31$-0.35 (that were almost absent in our cases; Figure~\ref{fig:ye}, bottom) would overproduce the nuclei of $A \sim 100$ (Figure~\ref{fig:nuclei}, top) that had, however, little effect on radioactive heating (Figure~\ref{fig:qdot}, bottom). Similarly, additional ejecta with $Y_\mathrm{e} > 0.45$ might have little effect on heating (except for the decay chain from $^{56}$Ni at late times).}. Different from a previous thought, $\beta$-decay contribution from $A\sim 130$ is unimportant in this case (Figure~\ref{fig:qdot_isotope}, top left). Our result (mFE-a) is consistent with additional contributions from fission and $\alpha$-decay (Figure~\ref{fig:luminosity}, middle left) but only at a qualitative level when considering uncertainties in nuclear ingredients. Currently, therefore, only the inferred $X_\mathrm{l}$ gives a hint to the amount of $r$-process elements. Our result in mFE-a, $X_\mathrm{l} = 0.035$, is a few times greater than the upper bound of the literature values \citep[$\approx 10^{-4}$--$10^{-2}$, e.g.,][]{Arcavi2017, Chornock2017, Nicholl2017, Waxman2017}. If we took it literally, the $r$-process mass of $\approx 0.01\, M_\odot$ in mFE-a would be regarded as the upper limit. The true value may be a few times smaller, which is still consistent with the theoretical range of dynamical ejecta masses in a recent work \citep[$= 0.002$--0.016,][]{Shibata2017}. However, no information can be obtained from the inferred $X_\mathrm{l}$ alone on the prodution of heavy $r$-process elements such as gold and uranium. Only a viable strategy appears to search a signature of spontaneous fission (from $^{254}$Cf and possibly a few Fm isotopes), which can affect late-time luminosities \citep[][]{Wanajo2014, Hotokezaka2016, Zhu2018}. As described in \S~\ref{sec:heating}, many uncertainties are involved in estimating the fission contribution. Nevertheless, a variation of late-time light curves among future NS merger events will be indicative of actinide production (Figure~\ref{fig:luminosity}, grey area), which is expected from spectroscopic studies of Galactic halo stars. In this paper comparison was made only with the NS merger GW170817, the first and currently unique detection of such an event. It should be noted that the same heating rates of our model (mFE-a) may not necessarily be applicable to future NS merger events. The reason is that the abundance distribution of light trans-elements may not be robust, which is very sensitive to $S$ and $Y_\mathrm{e}$ \citep{Meyer1998, Wanajo2018}. Small shifts of these quantities in post-merger disk outflows would substantially modify the abundance pattern. As can be seen in Figure~\ref{fig:qdot} (bottom), the behaviors of heating rates are very different among different $Y_\mathrm{e}$ groups for $Y_\mathrm{e} > 0.30$. The amounts of ejecta mass from disk outflows can also be dependent of, e.g., the masses and their ratio of merging NSs \citep[e.g.,][]{Shibata2017}. Even so, it is encouraging that we have means of discriminating between light trans-iron and $r$-process dominant mergers through the light curves of kilonovae in the future. | 18 | 8 | 1808.03763 |
1808 | 1808.03249_arXiv.txt | Neutrino induced reactions are a basic ingredient in astrophysical processes like star evolution. The existence of neutrino oscillations affects the rate of nuclear electroweak decays which participates in the chain of events that determines the fate of the star. Among the processes of interest, the production of heavy elements in core-collapse supernovae is strongly dependent upon neutrino properties, like the mixing between different species of neutrinos. In this work we study the effects of neutrino oscillations upon the electron fraction as a function of the neutrino mixing parameters, for two schemes: the $1+1$ scheme (one active neutrino and one sterile neutrino) and the $2+1$ scheme (two active neutrinos and one sterile neutrino). We have performed this analysis considering a core-collapse supernovae and determined the physical conditions needed to activate the nuclear reaction chains involved in the r-process. We found that the interactions of the neutrinos with matter and among themselves and the initial amount of sterile neutrinos in the neutrino-sphere might change the electron fraction, therefore affecting the onset of the r-process. We have set constrains on the active-sterile neutrino mixing parameters. They are the square-mass-difference $\Delta m^2_{14} $, the mixing angle $\sin^2 2\theta_{14}$, and the hindrance factor $\xi_s $ for the occupation of sterile neutrinos. The calculations have been performed for different values of $X_{\alpha}$, which is the fraction of $\alpha$-particles. For $X_{\alpha}=0$ the r-process is taking place if $\Delta m^2_{14} \geq 2 \, {\rm eV}^2 $, $\sin^2 2\theta_{14} < 0.8 $ and $\xi_s < 0.5$. For larger values of $X_{\alpha}$ the region of parameters is strongly reduced. The present results are compared to results available in the literature. | \label{Intro} Several experiments, like LSND \cite{aguilar01}, SK \cite{sk98}, SNO \cite{sno01}, and MiniBooNE \cite{aguilar07} among others \cite{chooz09,sage09,kamland03,gno05,k2k06,borexino08}, have provided evidences of neutrino oscillations caused by non-zero neutrino masses. In particular, the results of LSND and MiniBooNE are compatible with the inclusion of at least an extra sterile neutrino \cite{athanassopoulos95,athanassopoulos96,aguilar07b}. The consequences of the inclusion of massive neutrinos and sterile neutrinos in different astrophysical scenarios have been examined recently \cite{boyarsky09,mohapatra04,raffelt01}. The presence of massive neutrinos affects the rates of nuclear reactions where they participate, therefore a reformulation of the weak decay rates in terms of neutrino oscillation parameters and couplings is needed to explain several astrophysical processes, such as Big Bang Nucleosynthesis \cite{civitarese14} and nuclear reaction chains in stellar media \cite{qian93,qian95}. In supernovae, near the stellar-core, the neutrino flux might suffer conversions to the sterile sector causing a reduction in the number of electron neutrinos \cite{molinari03}. The effects of the active-active and active-sterile oscillations in supernova explosions have been studied by several authors \cite{balasi15,fetter03,balantekin04,tamborra12,wu15,janka12,pastor02,woosley94,qian03}. It allows to analyse the behaviour of matter at high densities and test properties of neutrino physics \cite{wu15,fetter03}. The process responsible for the production of nuclei heavier than iron is the rapid neutron-capture process or r-process. To be effective the reaction chain needs a neutron-rich environment, i.e. an electron fraction per baryon lower than $0.5$, a sufficiently large entropy, and short reaction times. The neutrino-driven matter-outflow, with a time of post-bounce of the order of 10 seconds and high neutron density, is a favoured mechanism for the production of elements heavier than iron \cite{qian03}. As said before, due to neutrino oscillations, the neutron abundance of the wind is modified when sterile neutrinos are included \cite{qian96}. The neutrinos determine the neutron to proton ratio, thus a possible conversion between flavours could alter this rate, and consequently the conditions needed for the r-process to take place. Calculations performed in the context of semianalytical models have shown that it is difficult to achieve the generation of the required number of free neutrons \cite{hoffman97}. The inclusion of massive neutrinos in the formalism which describes supernovae explosions affects the cross sections involved in the reaction chains that produce heavy nuclei by modifying the abundances of the elements ejected into the interstellar medium and by changing the neutrino number densities \cite{mclaughlin99,tamborra12}. In this paper, we study the impact of neutrino oscillations upon the electron fraction in the late neutrino-driven wind epoch. This work is organized as follow. In Section \ref{estrella} we present a brief description of the supernova environment and the determination of the electron fraction. In Section \ref{densidades} we present the formalism needed to calculate the neutrino densities and show the results of the calculations in Section \ref{resultados}. In order to estimate the differences between our results and the results obtained by applying other approximations \cite{keil03,tamborra12,pllumbi15}, we have performed calculations of the electron fraction $Y_e$ as a function of the stellar radius, by using power-law and Fermi-Dirac distributions, as discussed in section \ref{resultados}. The conclusions are drawn in Section \ref{conclusiones}. | \label{conclusiones} In this work we have studied the impact of the inclusion of massive sterile neutrinos upon the physical conditions required for the occurrence of the r-process in supernovae. The analysis was performed by calculating the electron-fraction in the stellar interior as a function of the sterile-active neutrino mixing parameters. We have found that the electron abundance is sensitive to the inclusion of sterile neutrinos, and that it depends on the neutrino interactions, e. g. neutrino-neutrino interactions or neutrino-matter interactions. From our results, the onset of the r-process in presence of a sterile neutrino, is compatible with the limits $\Delta m_{14}^2 \geq 2 {\rm eV}^2$, $\sin^2 2\theta_{14}<0.8$, and $\xi_s <0.5$, for the square-mass-difference, mixing-angle and enhancement factor of the sterile neutrino sector, respectively. As explained in the text, we found that the r-process is strongly affected by the fraction $X_{\alpha}$. | 18 | 8 | 1808.03249 |
1808 | 1808.06977_arXiv.txt | We present a method that enables wide field ground-based telescopes to scan the sky for sub-second stellar variability. The method has operational and image processing components. The operational component is to take star trail images. Each trail serves as a light curve for its corresponding source and facilitates sub-exposure photometry. We train a deep neural network to identify stellar variability in wide-field star trail images. We use the Large Synoptic Survey Telescope (LSST) Photon Simulator to generate simulated star trail images and include transient bursts as a proxy for variability. The network identifies transient bursts on timescales down to 10 milliseconds. We argue that there are multiple fields of astrophysics that can be advanced by the unique combination of time resolution and observing throughput that our method offers. | \label{sec:intro} The universe remains relatively unexplored on short time scales in the optical region of the electromagnetic spectrum. Charge coupled devices (CCDs) have been the detectors of choice in astronomy for over four decades. These large pixel arrays provide high quantum efficiency with excellent spatial resolution, but are conventionally operated in integration mode with exposure durations of order tens of seconds. Such long exposures preclude these instruments from imaging astrophysics that manifests on shorter time scales. High-speed photometric surveys, which we define as imaging large areas of sky with time resolution below one second, have been mostly unavailable, both for the study of known variable sources and for the search for new phenomena. The violent and rapidly varying radiation from black holes, neutron stars, and white dwarfs makes them promising targets for high time resolution imaging. The rotation, pulsation, and local accretion dynamics of these compact stellar remnants tends to occur on time scales ranging from seconds to milliseconds. Their extreme density also makes them an excellent testing ground for nuclear, quantum, and gravitational physics \citep{2004Sci...304..536L, 2016PhRvD..94h4002Y, 2017Natur.551...80K}. Compact stellar remnants are one of many applications. High-speed optical photometry has also supplemented the study of brown dwarfs, cataclysmic variable stars, eclipsing binary stars, X-ray binary stars, extrasolar planets, flare stars, active galactic nuclei, asteroseismology, and atmospheres of solar system objects \citep{1988Natur.336..452H, 2007MNRAS.378..825D}. Combining high-speed imaging with a wide-field instrument opens up the possibility of serendipitously observing a Kuiper Belt occultation, the immediate afterglow of a gamma ray burst, an optical counterpart of a fast radio bursts, and other rare phenomena \citep{2013AJ....146...14Z}. The majority of existing high-speed optical imaging instruments, such as ULTRACAM, ULTRASPEC, CHIMERA, and HiPERCAM take advantage of frame transfer and electron multiplying CCDs \citep{2007MNRAS.378..825D,2014MNRAS.444.4009D,2016SPIE.9908E..0YD, 2016MNRAS.457.3036H}. By limiting the camera readout to a small window surrounding a source of interest, these instruments can achieve high sample rates, even over 1,000 Hz. In this work, we are primarily interested in developing the capacity to detect sub-second stellar variability over a \textit{wide field of view}. ULTRACAM for example, can image a $1024 \times 1024$ pixel array at 40 Hz. Another promising approach for scanning the sky for sub-second variability is to operate CCDs in a continuous readout mode. \cite{2009AJ....138..568B} achieved 200 Hz photometry with continuous readout on MEGACAM. We show how similar performance can be achieved at existing facilities through a relatively trivial modification to the observing plan. We revisit an idea originally introduced by \cite{1986PASP...98..802H}: that star trails provide sub-exposure time resolution. Almost every major telescope is capable of producing star trail images by simply turning off the tracking. The principle hurdle to this method is the need to process these unorthodox images. For this, we leverage deep learning. Deep learning has achieved impressive results in many areas of machine learning. While neural networks and other paradigms in the field have a long history \citep{McCullochPitts43,Hochreiter:1997:LSM:265493.264179,lecun-gradientbased-learning-applied-1998,Hinton:2006:FLA:1161603.1161605,Bengio:2009:LDA:1658423.1658424}, increases in computational capacity, data sizes, and a series of practical tricks \citep{Glorot10understandingthe,DBLP:journals/corr/abs-1207-0580,journals/corr/KingmaB14,Ioffe:2015:BNA:3045118.3045167} have allowed deep neural networks to recently realize their full potential. From 2012 to 2014 deep learning methods broke through stubborn hurdles in object recognition \citep{NIPS2012_4824}, natural language processing \citep{Sutskever:2014:SSL:2969033.2969173}, and speech recognition \citep{graves2013speech}. In 2014, Ian Goodfellow introduced the powerful paradigm of generative adversarial networks \citep{NIPS2014_5423}. Iconic successes such as reaching human level performance in Atari games \citep{mnih2013playing,mnih-dqn-2015} and beating Go professionals \citep{SilverHuangEtAl16nature,silver2017mastering} also brought significant attention to these methods. Astronomy is a data rich field and presents many opportunities for deep learning. \cite{2015MNRAS.450.1441D} used a convolutional neural network to predict galaxy morphologies in the Galaxy Zoo project. \cite{2017MNRAS.467L.110S} used a generative adversarial network to recover astrophysical features in images beyond the deconvolution limit. \cite{2018MNRAS.473.3895L} developed a method to find galaxy-galaxy strong lenses. \cite{2017Natur.548..555H} reduced the time to analyze strong gravitational lenses by seven orders of magnitude. \cite{2018AJ....155...94S} identified new exoplanets. \cite{2017arXiv171107966G} improved gravitational wave detection. \cite{2017arXiv170906257M} classified stellar light curves. Finally, \cite{2017arXiv171001422S} used deep learning for image subtraction, which is closest to our application. We design a deep neural network to sift through wide field star trail images and detect variability. The input to the network is a simulated star trail image and the output is an image containing only the excess flux from variability. We argue that this technique is well suited for the diverse aforementioned applications. This paper is organized as follows. In Section \ref{sec:data} we describe the various LSST simulators and how we use them to produce training, development, and test datasets for our network. In Section \ref{sec:network} we describe our network architecture and training process. In Section \ref{sec:results} we assess the performance of our network with qualitative and quantitative measures. In Section \ref{sec:discussion} we review avenues for future work. In Section \ref{sec:conclusion} we summarize our findings. | \label{sec:conclusion} To conclude, we summarize our main results as follows. \begin{enumerate} \item We train a deep neural network to detect bursts in simulated LSST star trail images. \item We demonstrate that our network is robust. The bursts the network fails to detect are intrinsically difficult cases. The network performs well on the test set and it generalizes to new types of variability. \item We empirically confirm that star trails enable sub-exposure time resolution. The combined network has less than 1\% false-positives, detects variability in sources out to 20th magnitude, and detects bursts on time scales down to 10 milliseconds. \end{enumerate} The primary implication of this work is that by taking star tail images and processing them with deep learning we can extend the scientific payload of large survey telescopes to the vast field of high time resolution astrophysics. A number of astrophysics communities can benefit from such observations. | 18 | 8 | 1808.06977 |
1808 | 1808.10493_arXiv.txt | {Since 2013, the highest energy IceCube cascade showers overcame the common muon neutrino tracks. This fast flavor changes, above few tens TeV, has been in debt to the injection of the long-searched astrophysical neutrino. However for what concern the recent published 54 neutrino High Energy Starting Events (HESE) in 2016, as well as the most recent ones of 82 and 103 IceCube events (2017-2018) and the several dozens of thorough-going muon tracks formed around the IceCube, none of them are pointing or clustering toward any expected x, gamma or radio sources: no one in connection to GRB, no toward active BL Lac, neither to AGN source in Fermi catalog. No clear correlation with nearby mass distribution (Local Group), nor with galactic plane. Withal, there have not been any record of the expected double bang due to the tau neutrino birth and decay among several events above 200~TeV energy (we are disregarding for a moment the most recent two tau possible identification); no any self-clustering events at tens TeV energy raised in most recent search. Furthermore, there is a tension between the internal HESE event spectra power index and the external thorough-going muon tracks one. As we will show at the conclusions a more mundane (but a bit more abundant) prompt charmed atmospheric neutrino component may pollute and rule the data, explaining most of the present enlisted IceCube puzzles. We review the last HESE event data shown in early and in most recent papers (and talks in Neutrino 2018) making the case for the simplest conclusions. A very recent, unique, celebrated thorough-going muon in IceCube on $22/09/2017$ track possibly correlated to a gamma AGN (TXS~0506+056), and two revisited HESE events of possible tau neutrino nature, are mitigating the IceCube tau absence, but as we show more abundant tau signals are needed to confirm an astrophysical nature. A better filtered and guaranteed neutrino astronomy it is required: the tau flavor ones above TeVs have negligible or none (oscillated atmospheric) polluted background; very few percent of highest atmospheric charmed events may contribute to the noise (possibly the two event observed by IceCube). Astrophysical tau are possibly arising at best at highest energies (PeVs ones) ultimately overcoming noises and their tracks at hundred PeV edges may overcome the same atmospheric muon tracks. Tau neutrino astronomy is therefore the best road to highest energy neutrino astronomy. It may be revealed in double bangs in future IceCube events, by their tau birth first and their later decay in ice. A much sharp signal is made by upgoing-horizontal tau air-showers originated within mountain chain or inside Earth crust and better observable from mountains, balloons or from space, by large size array detectors (PAO observatory, Telescope Array, MAGIC telescope, Ashra, GRAND, POEMMA, ANITA). Some of these tau events might be already hidden in most recent ANITA up-going records at tens or hundred PeV. In conclusion we claim that in IceCube events are probably still sunk in a very polluted conventional and mostly prompt neutrino noise. Their dominance it is probably due to atmospheric charm emergence at a little higher rate than the expected one. We believe therefore that most astrophysical neutrino signals are still hidden below the ashes of these new, anyway discovered, prompt atmospheric noise. } \begin{document} | \label{sec:intro} Since November 2013 the sudden change of Ultra High Energy (UHE) neutrino flavor (IceCube HESE) was discovered by IceCube; its revolution has been discussed \cite{Fargion_2014} in the framework of the long-awaited birth of a UHE$\nu$ astronomy \cite{2017arXiv170200021F, 2017arXiv170500383K, 2015EPJWC..9908002F, 2014NuPhS.256..213F, FARGION2008146, 1742-6596-110-6-062008}. Indeed, up to TeV energy the UHE$\nu$ are following the expected spectra of atmospheric $\nu$ secondaries for which it is found a flux ratio of $\nu_\mu$ ($\bar{\nu}_\mu$) over $\nu_e$ ($\bar{\nu}_e$) of $\phi_{\nu_\mu}:\phi_{\nu_e}\approx20:1$. Suddenly, at few tens TeVs, the flavor revolution raised: four times more shower or cascade events, started by $\nu_e$, or by $\nu_\tau$ or by Neutral Current (NC) events, have shown respect the muon tracks events in IceCube: $\phi_{\nu_\mu}^{\text{tracks}}:(\phi_{\nu_e}+\phi_{\nu_\tau}+\phi_{\text{NC}})=1:4$ \cite{2013arXiv1309.7003I}. Since 2015 IceCube declared its ability to reveal double shower by tau at highest ($> 200$) TeV energies. Among the highest ten events above 200~TeV energy, no tau double cascades was detected. Indeed, above 200~TeV neutrino energy the so-called double chain of events, producing first a shower ($\nu_\tau$+N) and then a consequent $\tau$ with its forward decay (ten meter or more) in a second shower (double bang mechanism) \cite{1995APh.....3..267L}, was in principle observable but experimentally still undetected not in 2013 nor in 2015 \cite{2016PhRvD..93b2001A}, neither it was in more recent analysis \cite{2017arXiv170205238X}. In point of fact, at least a 10 of such events above 200~TeV were not revealing any $\tau$. Now, the $\nu_\mu$ and $\nu_e$, as known, are extremely polluted by atmospheric noises, while $\nu_\tau$ (and $\bar{\nu}_\tau$) are the unique beyond-doubt probe of an extraterrestrial signal. Therefore we face the urgent question: why tau has not been detected \cite{Fargion_WHY} or better to say it is so rare? \subsection{The article structure}\label{sec:intro_1} The article is structured as follows: in next section~\ref{sec:2} we remind the missing element in the desired $\nu$ astronomy, then we briefly review in section~\ref{sec:3} The UHE$\nu$ astronomy encoded in last 103 IceCube event: the last 54, 82 and 103 UHE neutrino events facing that key question. We than recall in subsection~\ref{sec:3_1} the effective volume and area in each neutrino flavor detection. Therefore we discuss in section\ref{sec:4} The recent 83 HESE and last 103 HESE IceCube events; in subsection~\ref{sec:4_1} the tau expected and the observed rates. These arguments required a remind on the effective volume and area of detection for each neutrino flavor. We had to underline our recent first surprise in subsection~\ref{sec:4_2}: double HESE sample (from 54 to 103) with less events above 200 TeV. This metamorphosis in IceCube catalog is much more evident in last 103 more radical events displacements. Therefore we show in section~\ref{sec:5} how to estimate the astrophysical tau neutrino presence in 54 or 103 sample. We introduce in subsections the two very recent claim of a tau signal among UHE IceCube events; in subsection~\ref{sec:5_1} a weighted rate probability for each neutrino flavor is discussed, in particular we concentrated on the last 36 IceCube events above 100 TeV that might be observed by double bang or pulse by last improved analysis: in subsection~\ref{sec:5_2} we discuss a weighted rate probability for each neutrino flavor. In following subsection~\ref{sec:5_3} the most recent $36$ UHE$\nu$ events above 100 TeV, candidate to be disentangled by eventual double bang, they were considered with their statistical weights for each flavor in IceCube. On the largest sample of $36$ UHE$\nu$ detectable we restricted to the 28 (only shower events) events in estimating the probability to observe a tau signal. In following subsection~\ref{sec:5_4} we considered in more details the largest sample of UHE neutrino detectable tau sample above 100 TeV threshold. In connection with the secondary tau neutrino role in charmed atmospheric events we added a subsection~\ref{sec:5_5} on the charmed neutrino flux and their tau minor presence. Their charmed foreseen events are in reasonable agreement with none, one or two tau signals. The conclusion are summarizing the results in section~\ref{sec:conclusion} , see Fig. \ref{Fig18}. The article is followed by a brief appendix with three figures that are reminding: $1)$, the tension on astrophysical spectra model for HESE and for the trough going muons, see Fig.\ref{Fig11}; $2)$ the average flavor balance that (up to last year) was derived by IceCube, (for electron, muon and tau component) at ratio $1,1,0$m see Fig.\ref{Fig12} (as for the charmed atmospheric case prescribed) ; $3)$ in the appendix last figure, see Fig. \ref{Fig15} we remind the remarkable Glashow resonance peak role in $\nu$ cross sections: its presence leading to $6.3$ PeV shower (or double bang) it is still unobserved. On the contrary the Glashow resonance it would be probably already taken place under a hard astrophysical spectra, while it might be still compatible, being hidden within a charmed atmospheric and softer spectra, as our thesis suggest. | \label{sec:conclusion} The main message of present article it is the following: there are already many hints and strong statistical signals that the IceCube events might be mainly polluted by atmospheric prompt neutrinos and very marginally by astrophysical ones. The polluted atmospheric events are charmed ones at a rate just a little above the foreseen ones \cite{2008Reno_TAU_PROMPT}. This may be indebt to the proton (and not other nuclei) dominance at PeVs (few, tens, hundred) energies. The probability that the UHE neutrino are astrophysical or charmed for a given number of tau detection it is summarized in figure~\ref{Fig18} where the probability that just two tau were observed as astrophysical ones is quite negligible ($<99.5\%$), and it is at least 50 times less probable than an atmospheric charmed signature. Naturally IceCube results are still preliminary and consequently our conclusions are as waiting for more definitive data. However there are several arguments, shown in our article, suggesting already the need for a radical reinterpretation of HESE (eventual) astrophysical nature. | 18 | 8 | 1808.10493 |
1808 | 1808.07971_arXiv.txt | Sensor pattern noise has been found to be a reliable tool for providing information relating to the provenance of an image. Conventionally sensor pattern noise is modelled as a mutual interaction of pixel non-uniformity noise and dark current. By using a wavelet denoising filter it is possible to isolate a unique signal within a sensor caused by the way the silicon reacts non-uniformly to light. This signal is often referred to as a fingerprint. To obtain the estimate of this photo response non-uniformity multiple sample images are averaged and filtered to derive a noise residue. This process and model, while useful at providing insight into an images provenance, fails to take into account additional sources of noise that are obtained during this process. These other sources of noise include digital processing artefacts collectively known as camera noise, image compression artefacts, lens artefacts, and image content. By analysing the diversity of sources of noise remaining within the noise residue, we show that further insight is possible within a unified sensor pattern noise concept which opens the field to approaches for obtaining fingerprints utilising fewer resources with comparable performance to existing methods. | Sensor pattern noise (SPN) is a reliable tool for tracking the provenance of images \cite{lukas2006digital}. Through the use of high-pass filtering, a unique signal can be extracted from an image consisting of Photo-Response non-uniformity (PRNU) noise. This signal is unique to the image sensor and is capable of discrimination across cameras of the same make and model. This discrimination is because the PRNU is defined as the pixel to pixel variance in output intensity of an image sensor when illuminated with a constant light source. The PRNU is the light-sensitive signal caused by Pixel Non-Uniformity (PNU) within a discrete image sensor. It is statistically unlikely for two image sensors to have the same PRNU fingerprint. This capability has been demonstrated with a false acceptance rate of 0.0024\% and a false rejection rate of 2.4\% making PRNU comparison an attractive tool where other evidence can also be used to verify the outcome \cite{goljan2009large}. The PRNU approach was further reinforced experimentally by \cite{farid2016photo} showing that a camera with a positive match to an image with the same PRNU is 1/100,000 or 99.999\%. Since images in the real world are often never illuminated with a constant light source, solving the blind source camera identification problem in this manner requires large sample sizes of images specially crafted to ensure scene contamination is minimised. Processing the large sample of images is either time consuming or requires large amounts of computing resources to ensure efficacy. These resources are not always available for forensic investigators in the field who need efficient tools to quickly and accurately quarantine evidence. There are further challenges in maintaining chain of evidence for embedded cameras, such as in contemporary smart phones and the emerging field of wearable technologies. Through careful analysis of the current sensor technology in terms of optical effects, semiconductor physics and the environment image sensors operate in, this paper considers the current methodology for measuring the unique PRNU signal and shows that other options exist for extracting the unique signal that may provide more accurate results and lead to the development of more efficient tools. | SPN has shown to be a promising area of research for answering provenance questions relating to imagery. It still suffers from reliance on large data sets ideally constructed from flat fielded images and a-priori information that is not always apparent or readily available to forensic investigators in the field. By rethinking the noise model of a digital camera, SPN can be isolated as an element that is primarily dependent on the physical silicon that each image sensor is built from, regardless of technology. From this analysis, it can be seen that there are alternative ways to create the unique fingerprint that will not rely on large data sets or mass computational resources. An example, the subject of current work, is the exploitation of dark current. | 18 | 8 | 1808.07971 |
1808 | 1808.05974_arXiv.txt | {Using high-resolution data of the $^{12}$CO and $^{13}$CO ({\COtrans}) line emission from The Mopra Southern Galactic Plane CO Survey in conjunction with neutral hydrogen observations from the Southern Galactic Plane Survey (SGPS) and mid-infrared {\Spitzer} data, we have explored the large-scale environment of the supernova remnant {\kes}. On the basis of these data, we identified for the first time the parent cloud of {\kes} in its whole extension and surveyed the HII regions, masers, and the population of massive young stellar objects in the cloud. The whole unveiled giant cloud, located at the kinematic distance of 12.0$\pm$3.6~kpc, whose average total mass and size are $\sim$10-30~$\times10^5$~{\Msun} and $\sim$$26^{\prime}$, also shines in $\gamma$-rays, as revealed by the Large Area Telescope on board the {\itshape Fermi} satellite. We determined a high average proton density $\sim$500-1000~\cm{-3} in the large molecular complex, of which protons from the neutral atomic and ionised gases comprise only $\sim$15\%.} | Massive stars as well as their stellar remnants created after supernova (SN) events spatially coexist in their parent molecular cloud (MC) \citep{mck07}. The accelerated particles created at the SN shock fronts and/or in star-forming regions (SFRs) can cause emission at $\gamma$-ray energies as the result of their interaction with ambient photons and/or with the interstellar medium (ISM) in which these objects evolve \citep{banik+17, rom10}. Knowledge of the physical conditions of the ambient matter in bright $\gamma$-ray emitting molecular complexes is a particular important key in order to obtain information on the creation and distribution of Galactic cosmic rays. In these regards, surveys designed to cover the large-scale distribution of different molecular species and of atomic hydrogen \citep{hey15, kal09} constitute an ideal framework to assess these matters. We here present a revisited analysis of the environs of the supernova remnant (SNR) {\kes}, where an excess in $\gamma$ rays at GeV energies was detected with the {\itshape Fermi}-Large Area Telescope (LAT) \citep{liu+15}. The interaction of {\kes} with nearby molecular gas has been known since the 1990s with the discovery of OH-1720~MHz maser emission \citep{koralesky+98}. A small portion of this gas about $\sim$~8$^{\prime}$ in size that also has a radial velocity of $\sim$$-$50~km~s$^{-1}$ was investigated by \citet{zhang+15}. Now, based on the interesting results reported by these authors, we have begun the analysis of the ISM in a significatively larger region using new information of the molecular $^{12}$CO and $^{13}$CO lines emission from The Mopra Southern Galactic Plane CO Survey, as well as neutral hydrogen and mid-infrared data from the Southern Galactic Plane Survey (SGPS) and {\Spitzer}, respectively. From our study we identified for the first time the giant parent molecular complex where {\kes}, a number of massive young stellar objects, and star-forming regions are mixed. The paper is organised as follows: Sect.~\ref{data} describes the observations, and the complete characterisation of the discovered cloud along with the proton content determination are presented in Sect.~\ref{counterparts}. Sect.~\ref{summary} summarises our findings. The implications of the current study on the production of the observed $\gamma$ rays along with the modelling of the broadband emission from radio to $\gamma$-ray energies to understand the relative contribution of leptons and hadrons in the {\kes} region will be addressed in a companion paper (Supan et al. 2018b). | \label{summary} Using the high-quality data acquired as part of The CO Mopra Southern Galactic Plane CO Survey, we investigated in detail the physical properties of the $^{12}$CO and $^{13}$CO ({\COtrans}) gas emission in the direction to SNR~{\kes}. The molecular gas information presented in this work, used in conjunction with neutral atomic hydrogen observations from SGPS and {\Spitzer} mid-infrared data to describe the thermal gas, as well as HII, masers, and young massive stellar objects, provide a complete characterization of the SNR environment. Taking advantage of the large area covered by the observations, we uncovered the natal cloud of these objects, which is located at 12.0$\pm$3.6~kpc from us, with a size of $\sim$28$^{\prime} \times$18$^{\prime}$ and covers a broad velocity range from $\sim$$-$71 to $-$41~km~s$^{-1}$. This cloud matches (on the plane of the sky) the $\gamma$-ray emission detected by {\Fermi}. Compared with the previous work of \citet{zhang+15} in a much smaller region only around SNR~{\kes}, this is the first time that the molecular gas towards the $\gamma$-ray emission is analysed in its whole extent. For the large cloud, we found a total interstellar proton density of 460$\pm$160~cm$^{-3}$, while for the smaller region enclosing the $\gamma$-ray peak (within the relatively low angular resolution of {\Fermi}) and the molecular material adjacent to {\kes}, the proton density is $950\pm330$~cm$^{-3}$. Both estimates include contributions from the molecular, atomic, and ionised gases in the region. This work clearly demonstrates the effectiveness of the CO Mopra Survey in the quest of finding large interstellar complexes associated with $\gamma$-ray radiation, which provides high sensitivity and high spatial and spectral resolution. The case discussed here may add to a short list of the few well-known SNRs \citep[e.g. W28, W51C, and IC~443,][]{gab15} that interact with molecular material in massive SFRs. The implications of our analysis on the production of the $\gamma$-ray flux will be investigated in a separate paper (Supan et al. 2018b). | 18 | 8 | 1808.05974 |
1808 | 1808.02286_arXiv.txt | Prominences are incredibly dynamic across the whole range of their observable spatial scales, with observations revealing gravity-driven fluid instabilities, waves, and turbulence. With all these complex motions, it would be expected that instabilities driven by shear in the internal fluid motions would develop. However, evidence of these have been lacking. Here we present the discovery in a prominence, using observations from the Interface Region Imaging Spectrograph (IRIS), of a shear flow instability, {the Kelvin-Helmholtz sinusoidal-mode of a fluid channel}, driven by {flows} in the prominence body. This finding presents a new mechanism through which we can create turbulent motions from the flows observed in quiescent prominences. The observation of this instability in a prominence highlights their great value as a laboratory for understanding the complex interplay between magnetic fields and fluid flows that play a crucial role in a vast range of astrophysical systems. | Solar prominences are cool plasma suspended in the $10^6$\,K solar corona by magnetic fields \citep{TAN1995}. Space-based observations has revolutionised our understanding of prominences, where we now know that they are incredibly dynamic across the whole range of observable spatial scales \citep{MAC2010}. Investigations show that the dynamical features observed in prominences both drive and are driven by gravity-driven fluid instabilities \citep{BERG2010, BERG2011, HILL2011b, HILL2012, HILL2018}, waves \citep{ARRE2012, HILL2013, ANT2015}, and turbulence \citep{LEO2012, FREED2016, HILL2017}. The complex motions observed in prominences can be clearly seen to create shear flows, and so it would be expected that instabilities driven by this shear would develop. The classic shear flow instability is the Kelvin-Helmholtz instability (KHi), which breaks up coherent sheets of vorticity into vortices. This instability comes in two distinct flavours: the surface mode of the instability that drives vortex formation at the boundary between two non-parallel flows\citep{CHAN1961}, and {the modes that act on channel flows including the sinusoidal-mode which drives the development of serpentine patterns} \citep{DRAZINREID1981}. {Magnetic fields work to suppress the instability. For an arbitrary shear flow, stability of the flow is guaranteed unless the difference between the maximum and minimum velocities is twice the minimum Alfv\'{e}n velocity in the direction of the flow \citep{HUGHES2001}.} {The surface mode of the KHi has been observed in many astrophysical systems. This includes where the solar wind interacts with the flanks of the magnetosphere \citep[e.g.][]{HASE2004}, associated with erupting regions \citep{OFMAN2011}, on the flanks of coronal mass ejections \citep{FOU2011, MOSTL2013} and where emerging magnetic flux interacts with prominences \citep[e.g.][]{BERG2010, RYU2010, BERG2017}. The sinusoidal-mode of a channel flow has proved more elusive, but it is believed to be important in coronal plumes \citep{ANDR2001}, and astrophysical jets \citep{FERR1981}.} {There has been a wide range of numerical and analytical studies investigating the role of this instability in astrophysical settings, often in the context of explaining observations \citep[e.g][]{FOU2011, OFMAN2011, MOSTL2013}. \citet{MIURA1982} investigated the linear growth rate of the magnetic KHi for continuous compressible flows finding that increases in the width of the shear layer reduces the growth of the instability and that the instability can be suppressed by compressibility. % One important role of the nonlinear evolution of the KHi is its ability to develop turbulent flows through reconnection and secondary instabilities \citep[e.g.][]{MATSUMOTO2004}. This process has been seen in numerical studies of kink waves in the solar atmosphere, which found that the KHi can grow and become turbulent on the surface of oscillating coronal loops \citep[e.g.][]{TERR2008} and prominence threads \citep[e.g.][]{ANT2015}.} {There have been observations in the solar atmosphere, including in prominences, of the surface mode of the magnetic KHi % but observations of the {sinusoidal-mode} of the instability are still lacking. % Here we present the discovery in a prominence, using observations from the Interface Region Imaging Spectrograph \citep[IRIS;][]{DEPON2014}, of streams of fluid developing serpentine patterns as a result of becoming unstable to the KHi. } | One question that needs to be addressed is: in highly dynamic prominences, why do we not see this instability developing everywhere? In some cases it will just be that the angles of the instability to the line of sight are such that nothing can be observed even though the instability is growing. However, one part of this answer is likely to be that the magnetic fields are strong enough to at least delay the growth of the instability so that it does not become noticeable in many places. It is quite likely that, now we have discovered the presence of the KHi associated with prominence flows, many more instances will come to light. % {In these observations we found a shear flow instability for a velocity differences {between $\sim 15$ and $35$\,km\,s$^{-1}$} {and for this to be unstable we require $M_{\rm A}\gtrsim 2$.} As this condition does not include other important suppression mechanisms like continuous distributions, compression and viscosity (though the large Reynolds number will make this relatively unimportant), we can expect that the actual requirement for instability will be greater. Therefore, to achieve the fast development the instability as presented here velocity differences somewhere beyond this limit will be required.} An interesting result is that the flows presented here cannot be moving along magnetic field lines, unless the magnetic field strengths of this prominence are weaker than measured values \citep[e.g.][]{LEROY1989, CASINI2009, OROZCO2014, LEVENS2016}, but are moving almost perpendicular to the field, carrying it along with the flow. The high frequency with which small-scale flows appear means that the prominence magnetic field must be constantly twisted up and redistributed by these flows. Over the lifetime of a prominence this may lead to magnetic energy being transferred around the prominence, slowly evolving the global structure until it becomes unstable and erupts. {In our modelling we have not taken into account the optically thick prominence \ion{Mg}{2} emission. Though the characteristic speeds used in the modelling in this paper can be easily obtained by tracking prominence motions and the size of the flows can be measured, it is much harder to accurately determine the velocity and density distributions. Therefore, more work is necessary to go beyond the proof-of-concept simulations presented in this paper and to make it to possible to use this instability to directly infer the plasma and magnetic field conditions inside the prominence.} The discovery of this instability provides an explanation for how observed vortical motions \citep{LIGG1984} can be formed in a prominence and also has great implications for understanding the development of turbulence in prominences. % Investigations into this turbulence have revealed that a characteristic length scale of a few thousands of kilometres exists in the turbulence \citep{LEO2012, HILL2017}, where this length is a factor of a few longer in the vertical direction than the horizontal \citep{HILL2017}. The unstable flows presented here are at just the right size to drive turbulence from this scale inside the body of the prominence. The magnetic KHi is one of the fundamental instabilities of fluid dynamics, and this means that it regularly occurs % in many different astrophysical systems across a huge range of scales. Due to the commonality of the physics, by investigating this instability in one system we can learn about how it works in a wide range of other astrophysical systems. The high temporal and spatial resolution that is given by space-based, and will be given by future ground-based, observations of prominences means that we have an exceptional opportunity to investigate this important astrophysical phenomenon. | 18 | 8 | 1808.02286 |
1808 | 1808.00283_arXiv.txt | We present radio observations and modelling of one of the nearest and brightest Type IIP supernova SN\,2004dj exploded in the galaxy NGC 2403 at a distance of $\sim$ 3.5 Mpc. Our observations span a wide frequency and temporal range of 0.24 - 43 GHz and $\sim 1$ day to 12 years since the discovery. We model the radio light curves and spectra with the synchrotron emission. We estimate the mass-loss rate of the progenitor star to be $\dot{M}$ $\sim$ 1 $\times$ 10$^{-6}$ M$_{\odot}\, \rm yr^{-1}$ for a wind speed of 10 km\,s$^{-1}$. We calculate the radio spectral indices using 1.06, 1.40, 5.00 and 8.46 GHz flux density measurements at multiple epochs. We witness steepening in the spectral index values for an extended period predominantly at higher frequencies. We explain this as a signature of electron cooling happening at the supernova shock in the plateau phase of the supernova. We estimate the cooling timescales for inverse Compton cooling and synchrotron cooling and find that inverse Compton cooling is the dominant cooling process. | \label{sec:intro} Massive stars (M $>$ 8M$_{\odot}$) end their lives in spectacular explosions called core-collapse supernovae (SNe). Type II SNe are explosions of massive stars that retain their hydrogen envelope at the time of explosions and show copious hydrogen emission lines in their optical spectra \citep{filippenko1997}. Type IIP SNe is a sub-class of Type II SNe characterized by a pronounced plateau of nearly constant luminosity in the optical light curve following a maximum that lasts for $80-120$ days post explosion \citep{nadyozhin2003}. The plateau in their optical light curves is attributed to an extended hydrogen envelope intact to the star during explosion \citep{grasberg1971,falk1977,smartt2009b}. In a volume limited sample ($<$ 60 Mpc) of all core-collapse SNe, 48\% is comprised by SNe IIP \citep{smith2011}. Thus, SNe IIP is the most commonly observed variety of core-collapse SNe in the local universe and hence likely to be the most common evolutionary path in the end stages of the life of massive stars. Several lines of evidence from stellar evolution models and supernova (SN) light curve models suggest that the progenitors of SNe IIP are red-supergiants \citep[RSGs;][]{chevalier2006b}. The pre-supernova radius estimated from the plateau brightness and duration also falls in the typical RSG stellar radius \citep[$10^{2-3} R_{\odot}$;][]{grasberg1971,falk1977}. RSGs are also identified as the progenitor stars of SNe IIP SNe in direct detection efforts \citep{smartt2009b}. However, the masses of progenitor stars detected from pre-explosion images range from $8-15 M_{\odot}$ \citep{vandyk2003,li2005,li2006,maund2005a,maund2005b} which is closer to the lower mass end for a core-collapse event. No RSG star of mass $M > 17 M_{\odot}$ has been identified as a progenitor star of SNe IIP in direct detection efforts carried out in a volume limited sample \citep{smartt2009a}. In this context, a detailed census of type IIP SNe progenitors and the diversity of the properties of the SNe are important. In a Type IIP SN, the fast moving stellar ejecta interacts with the circumstellar medium (CSM) created by the stellar wind of the progenitor star. The interaction creates a strong shock that moves ahead of the ejecta and is called the forward shock \citep{chevalier1982}. Electrons are accelerated to relativistic energies at the forward shock and emit at radio frequencies. Radio emission is absorbed at early times by either free-free absorption (FFA) or synchrotron self absorption (SSA) and can be modelled as synchrotron emission affected by either or both absorption processes. Depending on the dominant absorption process, various physical parameters like mass-loss rate of the progenitor star, density of the CSM, ejecta density profile, shock deceleration parameter, magnetic field strength etc can be constrained from the modelled radio light curves and spectra \citep{chevalier1982,chevalier1998}. On the other hand, X-ray emission from Type IIP SNe can be either thermal or non-thermal in origin. Thermal X-rays can be emitted from the hot forward and/or reverse shock regions whereas non-thermal X-ray emission can be due to the Inverse Compton (IC) scattering. In Type IIP SNe, the plateau of the optical light curve is a phase of high density of ambient photons which can be IC scattered to X-ray energies by the relativistic electrons. As a result of IC scattering the relativistic electrons lose energy and there is a corresponding cooling break in the radio spectra above a characteristic frequency \citep{chevalier2006b}. In this work, we present long-term ($\sim$ 12 years) radio monitoring of a Type IIP, SN\,2004dj over a frequency range of 0.24 to 43 GHz. We model the radio light curves and spectra with the standard mini-shell model \citep{chevalier1982,chevalier1998} and derive the mass-loss rate of the progenitor star. We also search for the signatures of cooling in the radio spectra and interpret our results in light of other published results. The paper is organised as follows: In \S \ref{sec:sn2004dj-literature-review}, we discuss the previous work published on SN\,2004dj. In \S \ref{sec:obs}, we present the GMRT and VLA observations of SN\,2004dj in detail. In \S \ref{sec:modelling}, we describe the standard model for radio emission. In \S \ref{mass-loss-rate calculation}, we calculate the mass-loss rate of the progenitor star from the modelled parameters. In \S \ref{sec:cooling}, we discuss the evolution of spectral indices and the signatures of cooling in the radio spectra. We summarise our main results in \S \ref{sec:summary}. | \label{sec:cooling} \subsection{Progenitor properties} We model the radio observations with the standard mini-shell model \citep{chevalier1982} as explained in \S \ref{sec:modelling}. We estimate the shock deceleration parameter $m$ $\sim$ 0.9 (where $R \propto t^m$) for FFA model indicative of a mildly decelerating blast wave. This is in accord with the physics of the process as the blast wave is interacting with the CSM and is expected to slow down. The range of m values for Type II SNe is $m = 0.8 - 1$ \citep{weiler1986}. For a shocked shell of radius $R$ $\sim$ $t^{0.9}$, the ejecta density index n ($\rho \sim r^{-n}$) is given by $m = (n-3)/(n-2)$, assuming that the CSM is created by a steady stellar wind of density ($\rho \sim r^{-2}$) \citep{chevalier1982}. Thus for $m$ $\sim$ 0.9, we derive the power law ejecta density as $\rho \sim r^{-11.4}$. The progenitors of Type IIP SNe are understood to be RSGs with most of its hydrogen envelope intact during the SN explosion \citep{smartt2009b}. The value of $n$ for such a star is expected to be in the range $n = 7-12$ \citep{chevalier1982}. Thus the $n$ value derived from our analysis is consistent with a RSG progenitor of SN\,2004dj. Assuming FFA as the dominant absorption process, we derive the mass-loss rate of the progenitor star of SN\,2004dj as $\dot{M} = (1.37 \pm 0.11) \times 10^{-6}\, M_{\odot}\rm yr^{-1}$. The progenitors of Type IIP SNe are understood to be red-supergiants whose initial masses range from 8 - 25 $M_{\odot}$ \citep{heger2003}. For the best studied RSGs, the range of mass-loss rates is considerably large, i.e 2 $\times$ 10$^{-7}$ to 1.5 $\times$ 10$^{-5}$ $M_{\odot}\rm yr^{-1}$ \citep{jura1990,van2005}. The mass-loss rates of RSGs at the time of explosion is estimated using stellar evolutionary calculations as $\sim$ 3 $\times$ 10$^{-7}$ to 3 $\times$ 10$^{-5}$ $M_{\odot}\rm yr^{-1}$ \citep{schaller1992,chevalier2006b}. The theoretical estimate of mass-loss rate for 15$M_{\odot}$ models is derived as $\dot{M}$ = (0.84-1.6) $\times$ 10$^{-6}$ $M_{\odot}\rm yr^{-1}$ and for 20$M_{\odot}$ models is derived as $\dot{M}$ = (3.0-6.2) $\times$ 10$^{-6}$ $M_{\odot}\rm yr^{-1}$ \citep{chevalier2006b}. Thus the mass-loss rate derived for SN\,2004dj from radio observations are consistent with the theoretical predictions for a progenitor star of mass $\sim$ 15$M_{\odot}$. The progenitor mass of SN\,2004dj was infered as 15$M_{\odot}$ from stellar population studies of S96 cluster \citep{maiz2004a}. Mass-loss rate of the progenitor star of SN\,2004dj was derived by \cite{chakraborti2012} from X-ray emission measure as $\dot{M}$ = (3.2 $\pm$ 1.1) $\times$ 10$^{-7}$ $M_{\odot}\rm yr^{-1}$ which is $\sim$ 4 times lower than our mass-loss estimate. \cite{chugai2007} derived the mass-loss rate of SN\,2004dj to be $\sim$ 1 $\times$ 10$^{-6}$ $M_{\odot}\rm yr^{-1}$ for a wind velocity of 10 km\,s$^{-1}$, consistent with the mass-loss rate estimate from our analysis. \cite{chevalier2006b} derives a mass-loss rate of $\dot{M}_{-6}/w1$ $\sim$ (2-3) $T_{\rm cs5}^{3/4}$ for SN\,2004dj using the first 100 days of radio data. Here $\dot{M}_{-6}$ is the mass-loss rate in units of $10^{-6} M_{\odot}\rm yr^{-1}$ and $w1$ is the stellar wind velocity in units of 10 km\,s$^{-1}$. $T_{\rm cs5}$ denotes the CSM electron temperature in units of $10^{5}$ K. The authors compile the radio and X-ray data of all type IIP SNe available then and derive a range of mass-loss rates (1-10) $\times$ 10$^{-6}$ $M_{\odot}\rm yr^{-1}$ (see Table \ref{comparison-typeIIp}) for Type IIP progenitors and our mass-loss rate estimation is consistent with this range. \subsection{Signatures of cooling} \label{first-sub-section-cooling} The radio light curve and spectra of Type IIP SNe can be affected by cooling. Cooling becomes important depending on various parameters of SN such as ejecta, magnetic field strength, circumstellar medium, relativistic electrons etc \citep{chevalier2006b}. The electron can lose energy by adiabatic expansion, synchrotron cooling, and IC cooling. The dominant cooling process can be identified by calculating the break frequencies and cooling time scales for different cooling processes. One of the important signature of cooling is imprinted in the evolution of radio spectral indices. As a result of cooling the spectral index $\alpha$ steepens by $\Delta \alpha$ $\sim$ $-$0.5. This appears as a break in the spectra at a certain frequency when the electrons radiating above that characteristic frequency loses significant energy. Assuming that the synchrotron loss time scale is equal to the age of the SN, the expression for synchrotron break frequency is given by \cite{chevalier2006b} as \begin{eqnarray} \nu_{\rm syn} = 240 \left( \frac{\epsilon_{B}}{0.1} \right)^{(-3/2)} \left( \frac{\dot{M}}{10^{-6} M_{\odot} \rm yr^{-1}} \right)^{-3/2} \nonumber \\* \times \left( \frac{w}{10\, \rm km\,s^{-1}}\right)^{3/2} \left( \frac{t}{60\, \rm days} \right) \rm GHz \end{eqnarray} Assuming that the IC cooling time scale is equal to the age of the SN, the IC break frequency is \citep{chevalier2006b} \begin{eqnarray} \nu_{\rm IC} = 8 \left( \frac{\epsilon_{B}}{0.1} \right)^{(1/2)} \left( \frac{\dot{M}}{10^{-6} M_{\odot} \rm yr^{-1}} \right)^{1/2} \nonumber \\* \times \left( \frac{v_{w}}{10\, \rm km\,s^{-1}}\right)^{-1/2} \left( \frac{t}{60\, \rm days} \right) \nonumber \\* \left( \frac{V_{s}}{10^{4} \rm km\,s^{-1}}\right)^{4} \left( \frac{L_{\rm bol}(t)}{10^{42}\rm erg\,s^{-1}}\right) \rm GHz \end{eqnarray} \cite{chakraborti2012} derived $\epsilon_{B}$ = 0.082 and $\epsilon_{e}$ = 0.39 for SN\,2004dj using four epochs of \textit{Chandra} data ( $\sim$ 42, 56, 97 and 177 days post explosion) where the authors found prominent IC X-ray component in the first two epochs. We calculate the break frequencies corresponding to synchrotron cooling and IC cooling using the above equations at these two epochs (i.e day 42 and 56 post explosion). The value of $\epsilon_{B}$ is taken from \citep{chakraborti2012} and $\dot{M}$ from our calculations. We use the value of shock velocity $V_{s}$ = 9.2 $\times$ $10^{3}$ km\,s$^{-1}$ \citep{chakraborti2012}. The bolometric luminosity during the plateau is $L_{\rm bol}$ $\sim$ 0.89 $\times$ 10$^{42}$ ergs \citep{zhang2006}. Asuming a wind velocity of 10 km\,s$^{-1}$, and temperature of $10^{4}$ K, the $\nu_{\rm syn}$ and $\nu_{\rm IC}$ corresponding to day 42 and 56 days post explosion are $\sim$ 141, 4 GHz and 188, 5 GHz respectively. The synchrotron break frequency is too high to observe with the VLA during our observation epochs. The IC cooling break frequency is within our observation frequencies and with the multi-frequency observations of SN\,2004dj, we can look for this signature in the data. In Figure \ref{spectral-index variation}, We plot the evolution of spectral indices between successive frequencies, 1.06/1.39, 1.39/4.99, 4.99/8.46 GHz with time for SN\,2004dj. The spectral index values of 4.86/8.46 GHz approaches values $\sim$ $-$1 and lower during an extended period starting from $\sim$ day 50 (see Figure \ref{spectral-index variation}). Thus we see the IC cooling break at $\sim$ 5 GHz, roughly consistent with the above calculation. The optical bolometric light curve is in the plateau phase during the same period \citep{zhang2006}. The dense optical photon medium in the plateau phase of the SN provides seed photons and enhances the IC cooling. This kind of a dip in the spectral index is seen for the second time in a type IIP supernova after SN\,2012aw \citep{yadav2014}. We establish IC cooling as the dominant cooling process as there is evidence of cooling break at $\sim$ 5 GHz in the radio spectral evolution. Here, we calculate the cooling timescale at $\sim$ 5 GHz for IC and synchrotron cooling process to further investigate this. The ratio of synchrotron cooling time scale to the adiabatic expansion timescale is \citep{chevalier2006b} \begin{eqnarray} \frac{t_{\rm syn}}{t} \approx 2.0 \left( \frac{\epsilon_{B}}{0.1} \right)^{-3/4} \left( \frac{\dot{M}_{-6}}{v_{w1}}\right) ^{-3/4} \left(\frac{\nu}{10\, \rm GHz} \right)^{-1/2} \nonumber \\* \times \left( \frac{t}{10\, \rm days} \right)^{1/2} \end{eqnarray} The ratio of Compton cooling timescale to the adiabatic expansion timescale is \begin{eqnarray} \frac{t_{\rm Comp}}{t} \approx 0.18 \left( \frac{L_{bol}}{2 \times 10^{42} \, \rm erg\, s^{-1}} \right) ^{-1} \left( \frac{\epsilon_{B}}{0.1} \right)^{1/4} \left( \frac{\dot{M}_{-6}}{v_{w1}}\right) ^{1/4} \nonumber \\* \times V_{s4}^{2} \left(\frac{\nu}{10\, \rm GHz} \right)^{-1/2} \left( \frac{t}{10\, \rm days} \right)^{1/2} \end{eqnarray} In general, lower values of $\epsilon_{B}$ and $\dot{M}$ favours IC cooling and especially at early times. We estimate the ratios $t_{\rm Comp}/t$ and $t_{\rm syn}/t$ at $\sim$ 42 and 56 days post explosion using the above equations as $\sim$ 0.9, 5.3 and 1.1, 6.1 respectively. While the ratios of cooling time-scale is close to 1 for IC cooling, the ratios are much higher for synchrotron cooling; clearly indicates that synchrotron cooling is not plausible. Thus these numbers favour the IC cooling to be in operation during these epochs at the SN shock. This is in agreement with the detection of non-thermal IC component in the X-ray spectra of SN\,2004dj during the first two epochs of observations ($\sim$ day 42 and 56 post explosion) by \cite{chakraborti2012}. Thus both cooling time scale and break frequency calculations supports the IC cooling happening at the SN shock during the plateau phase. \begin{deluxetable*}{cccccccc} \centering \tablecaption{Comparison of SN\,2004dj parameters with other radio/X-ray bright Type IIP SNe. \label{comparison-typeIIp}} \tablecolumns{8} \tablenum{3} \tablewidth{0pt} \tablehead{ \colhead{SN} & \colhead{Parent} & \colhead{Distance} & $L_{\rm radio}$ & $L_{\rm X ray}$ & $\dot{M}$ & Progenitor mass & References \\ \colhead{-} & \colhead{Galaxy} & \colhead{(Mpc)} & (erg s$^{-1}$ Hz$^{-1}$) & (erg s$^{-1}$) & ($10^{-6} M_{\odot}\rm yr^{-1}$) & ($M_{\odot}$) & \\ } \startdata SN\,1999em & NGC\,1637 & 11.7 $\pm$ 1.0 & 2.2 $\times$ 10$^{25}$ & 9 $\times$ 10$^{37}$ & 0.9\tablenotemark{a} & $<$15 & 1,2,3,4 \\ SN\,1999gi & NGC\,3184 & 11.1$^{+2.0}_{-1.8}$ & --\tablenotemark{*} & 1.6 $\times$ 10$^{37}$ & $\sim$ 1\tablenotemark{b} & $<$12 & 1,2,3,5 \\ SN\,2002hh & NGC\,6946 & 5.5 $\pm$ 1.0 & 1.3 $\times$ 10$^{25}$& 4 $\times$ 10$^{38}$ & 1.2\tablenotemark{a} & -- & 3,6,7,8 \\ {\bf SN\,2004dj} & {\bf NGC\,2403} & {\bf 3.47} & {\bf 2.5 $\times$ 10$^{25}$} & {\bf 1.5 $\times$ 10$^{38}$} & {\bf 1.4\tablenotemark{c}}, (0.2 - 0.5)\tablenotemark{a}, 0.3\tablenotemark{b} & {\bf 15, $\sim$ 12, $>$20 } & 3,9,10,11 \\ SN\,2004et & NGC\,6946 & 5.5 $\pm$ 1.0 & 8.7 $\times$ 10$^{25}$ & (2.1 - 3)$\times$ 10$^{38}$ & (1.6-1.8)\tablenotemark{a}, $\sim$ 2\tablenotemark{b} & 15$^{+5}_{-2}$ & 3,6,12,13 \\ SN\,2011ja & NGC\,4945 & 3.36 $\pm$ 0.09 & 7.3 $\times$ 10$^{24}$ & (1.3 - 5.5) $\times$ 10$^{38}$ & 0.12 - 1.9\tablenotemark{b} & -- & 14,15 \\ SN\,2013ej & M74 & 9.6 $\pm$ 0.7 & --\tablenotemark{*} & $\sim$ (1 - 4) $\times$ 10$^{42}$ & 2.6\tablenotemark{b} & -- & 16,17 \\ SN\,2012aw & M95 & 10 & 7.2 $\times$ 10$^{25}$ & 9.2 $\times$ 10$^{38}$ & $\sim$ 3.2-10\tablenotemark{b} & 14-26, 17-18 & 18,19,20,21,22 \\ SN\,2017eaw & NGC\,6946 & 5.86 $\pm$ 0.76 & 1.7 $\times$ 10$^{25}$ & 1.1 $\times$ 10$^{39}$ & -- & -- & 22,23,24 \\ \enddata \tablecomments{This table is adapted from a similar table published by \cite{chevalier2006b} with more SNe added.} \tablecomments{References: (1) \cite{leonard2002}, (2) \cite{smartt2003}, (3) \cite{chevalier2006b}, (4) \cite{pooley2002}, (5) \cite{schlegel2001}, (6) \cite{li2005}, (7) \cite{pooley2002iauc}, (8) \cite{beswick2005iauc}, (9) \cite{chakraborti2012}, (10) \cite{maiz2004a}, (11) \cite{vinko2006}, (12) \cite{misra2007}, \cite{li2005}, (13) \cite{argo2005}, (14) \cite{chakraborti2013}, (15) \cite{mouhcine2005}, (16) \cite{chakraborti2016} (17) \cite{bose2015}, (18) \cite{kochanek2012}, (19) \cite{immler2012}, (20) \cite{fraser2012}, (21) \cite{vandyk2012}, (22) \cite{argo2017}, (23) \cite{bose2014}, (24) \cite{gref2017}.} \tablecomments{X-ray luminosity in the energy range of 0.5 - 8 KeV for all SNe except SN\,2011ja, SN\,2017eaw (0.3 - 10 KeV) and SN\,2012aw (0.2 - 10 KeV). Radio luminosity is spectral luminosity at $\sim$ 5 GHz. Where possible, the X-ray and radio luminosities are taken at the peak of the light curve and for others, the limits of peak flux density is used. The mass of the progenitor stars are from the pre-explosion observations of the SN site. The listed mass-loss rates are from X-ray/radio observations and modelling.} \tablenotetext{a}{The mass-loss rate is estimated from radio analysis and modelling assuming a stellar wind velocity of 10 km\,s$^{-1}$. While for SN\,1999em, SN\,2002hh, SN\,2004et \cite{chevalier2006b} and for SN\,2012aw \cite{yadav2014} calculate mass-loss rates assuming a CSM electron temperature of $10^{5}$ K, we recalculate these numbers for an electron temperature of $10^{4}$ K for comparison since we used electron temperature of $10^{4}$ K to calculate the mass-loss rate of SN\,2004dj.} \tablenotetext{*}{Not detected in radio.} \tablenotetext{b}{The mass-loss rates are estimated from the X-ray observations assuming a stellar wind velocity of 10 km\,s$^{-1}$.} \tablenotetext{c}{The mass-loss rate estimated for SN\,2004dj from this work assuming a stellar wind velocity of 10 km\,s$^{-1}$ and CSM electron temperature of $10^{4}$ K.} \end{deluxetable*} \cite{chevalier2006b} predicts a flattened light curve and a dip in the radio light curve of Type IIP SNe especially at higher frequencies as a signature of IC cooling. However, from our error bars on flux densities and cadence of observation, it is difficult to look for a flattening or small dip in the light curve as predicted by \cite{chevalier2006b}. Cooling processes and its effect on the radio light curves could be modelled better with a good quality and high cadence early time data at radio frequencies. \begin{figure*} \label{test} \includegraphics*[width=1\textwidth]{comparison.eps} \caption{ \scriptsize{Left panel:Radio spectral luminosity at $\sim$ 5 GHz of Type IIP SNe Right panel: X-ray luminosity in the energy range of 0.5 - 8 KeV for all SNe except SN\,2017eaw (0.3 - 10 KeV) and SN\,2012aw (0.2 - 10 KeV). Where possible, the X-ray and radio luminosities are taken at the peak of the light curve and for others, the limits of peak flux density is used. We include only the Type IIP SNe with both X-ray and radio detection.}} \end{figure*} \subsection{Circumstellar Electron Temperature} The circumstellar electron temperature is an uncertain parameter since the effect of SN radiation on the circumstellar gas needs to be estimated from the CSM models \citep{chevalier2003, chevalier2006b}. Calculations on Type IIL SNe suggest that the temperature at unit optical depth $T_{\rm cs}$ $\sim$ 3 $\times$ $10^{4}$ K for $\dot{M}_{-6}/w1$ = 3 and $T_{\rm cs}$ $\sim$ 1 $\times$ $10^{5}$ K for $\dot{M}_{-6}/w1$ = 10 \citep{lundqvist1988}. The mass-loss rate depends on $T_{e}$ as $\dot{M}$ $\propto$ $T_{e}^{0.68}$ \citep{weiler1986}, and hence the uncertainty in $T_{e}$ will cause significant errors in $\dot{M}$. Assuming an electron temperature of $T_{e}$ = 10$^{4}$ K, we derive the mass-loss rate of SN\,2004dj progenitor as $\dot{M} = (1.37 \pm 0.11) \times 10^{-6}\, M_{\odot}\rm yr^{-1}$ from our radio analysis and modelling. However, if we assume electron temperature $T_{e}$ = 10$^{5}$ K, the mass-loss rate will be $\dot{M} = (6.56 \pm 0.66) \times 10^{-6}\, M_{\odot}\rm yr^{-1}$. This is $\sim$ 20 times larger than the mass-loss rate derived from X-ray observations, i.e $(3.2 \pm 1.1) \times 10^{-7}\, M_{\odot}\rm yr^{-1}$ \citep{chakraborti2012}. It is also important to note that lower values of $\dot{M}$ favours IC cooling \citep{chevalier2006b,chakraborti2012}. If we use $\dot{M} = (6.56 \pm 0.66) \times 10^{-6}\, M_{\odot}\rm yr^{-1}$, the ratio of cooling timescales are $t_{\rm Comp}/t = 1.7 $ and $t_{\rm syn}/t = 1.9 $ on day 56 post explosion. Since both the ratios are much above 1, these numbers do not favour IC cooling. But there are several lines of evidences including the prominent IC X-ray component on day 56 post explosion \citep{chakraborti2012} showing that IC cooling is happening at the SN shock during this time as discussed in \S \ref{first-sub-section-cooling}. Thus the mass-loss rate derived for SN\,2004dj assuming a electron temperature of $T_{e}$ = $10^{5}$ K is too high to explain the observational signature of IC cooling. We suggest the electron temperature of the CSM to be $\sim$10$^{4}$ K rather than $10^{5}$ K for SN\,2004dj. We also note that \cite{misra2007} deduce the CSM temperature of another Type IIP SN\,2004et as 10$^{4}$ K from combined X-ray and radio data which has marginally higher wind density (see Table \ref{comparison-typeIIp}) and mass-loss rates \citep{misra2007} as that of SN\,2004dj. \subsection{Comparison of SN\,2004dj with other Type IIP SNe} Even though Type IIP is the most commonly observed variety of core-collapse SNe in the optical band, only few of them are known radio/X-ray emitters. SN\,2004dj being one of the best observed Type IIP SNe in the radio bands, here we compare the properties of SN\,2004dj with other radio/X-ray bright Type IIP SNe; SN\,1999em, SN\,1999gi, SN\,2002hh, SN\,2004et, SN\,2011ja, SN\,2013ej, SN\,2012aw and SN\,2017eaw. The physical parameters of these SNe along with the references are compiled in Table \ref{comparison-typeIIp}. We also plot the radio and X-ray luminosities of Type IIP SNe that was detected in both radio and X-ray bands in Figure 5. The radio and X-ray luminosities of SN\,2004dj is comparable to that of SN\,1999em and SN\,2002hh (see Fig 5) and is suggestive of comparable CSM densities. The mass-loss rate of the progenitor of SN\,1999em and SN\,2002hh are $\sim$ 0.9 $\times$ 10$^{-6}$ $M_{\odot}\rm yr^{-1}$ and $\sim$ 1.2 $\times$ 10$^{-6}$ $M_{\odot}\rm yr^{-1}$ respectively \citep{chevalier2006b} for a CSM temperature of $T_{e}$ = $10^{4}$ K. This is similar to the mass-loss rate derived for SN\,2004dj from our radio analysis (see Table \ref{comparison-typeIIp}). While the X-ray luminosity of SN\,2004et is comparable to that of SN\,2004dj, the radio luminosity of SN\,2004et is larger than ($\sim$ 3.5 times) SN\,2004dj and could be indicative of slightly larger wind densities. The mass-loss rate of SN\,2004et is derived from X-ray observations to be $\dot{M}$ $\sim$ 2 $\times$ 10$^{-6}$ $M_{\odot}\rm yr^{-1}$ \citep{misra2007} and $\dot{M}$ $\sim$ (1.6 - 1.8) $\times$ 10$^{-6}$ $M_{\odot}\rm yr^{-1}$ \citep{chevalier2006b}. These numbers are marginally larger than the mass-loss rate of SN\,2004dj derived from our analysis. SN\,2012aw has a larger X-ray luminosity ($\sim$ 6 times) and radio luminosity \citep[$\sim$ 3 times;][]{yadav2014} as compared to SN\,2004dj (see Fig 5). The mass-loss rate is $\dot{M}$ $\sim$ 3.2 $\times$ 10$^{-6}$ $M_{\odot}\rm yr^{-1}$ from X-ray analysis \citep{kochanek2012} which is larger than the mass-loss rate of SN\,2004dj. It is found that $\dot{M}$ depends on the metallicity of the regions of SNe as $\dot{M} \propto Z^{0.5}$ \citep{schaller1992,heger2003}. The metallicity of the regions of occurence of SN\,1999em \citep{smartt2003}, SN\,2004et \citep{li2005} and SN\,2004dj \citep{wang2005} are (1-2), (0.3-1) and $\sim$ 0.4 $Z_{\odot}$ respectively. This will have a minor effect in the $\dot{M}$ comparison discussed above. To summarize, the estimate of mass-loss rate of the progenitor star depends on various physical parameters including CSM electron temperature and metallicity of the cite of SN explosion. A fair comparison is not possible unless these quantities are well constrained by either observations or modelling. However, the mass-loss rates of Type IIP SNe progenitors deduced from radio and X-ray observations and modelling span over a range of $\sim$ (0.1 - 10)$\times$ 10$^{-6}$ $M_{\odot}\rm yr^{-1}$. We carried out detailed radio observations of Type IIP supernova SN\,2004dj at frequencies ranging from 0.24 - 43 GHz at ages from 1.12 days to $\sim$ 12 years post discovery. We model the radio observations with standard mini-shell model \citep{chevalier1982}. Both FFA and SSA models fit with the data reasonably well and it is difficult to conclude either of them as the dominant absorption process from the modelled parameters. However, from the optical line velocity measurements, we infer FFA as the dominant absorption process which is consistent with the prediction by \cite{chevalier2006b}. The radio observations and modelling are consistent with the interaction of the SN with an outer ejecta density profile $\rho$ $\sim$ $r^{-11.4}$ with a circumstellar density field created by a pre-SN steady stellar wind. We derive the shock deceleration parameter $m$ $\sim$ 0.9 ($R$ $\sim$ $t^{m}$) indicative of a mildly decelerating blast wave. Assuming FFA to be the dominant absorption process, we derive the mass-loss rate of the progenitor star as $\dot{M}$ = (1.37 $\pm$ 0.11) $\times$ 10$^{-6}$\, $M_{\odot}\rm yr^{-1}$. The mass-loss rate derived from our observations are consistent with the theoretical predictions for a progenitor star of mass $M \sim 15 M_{\odot}$. The mass-loss rate derived from our analysis is $\sim$ 3 times larger than the value derived by \cite{chevalier2006b} for SN\,2004dj from early radio data. However, our mass-loss estimate is consistent with the range of RSG mass-loss rates of type IIP SNe \citep{chevalier2006b}. \cite{chakraborti2012} estimated the mass-loss rate, $\dot{M}$ = 3.2 $\times$ 10$^{-7}$ $M_{\odot}\rm yr^{-1}$ from X-ray emission measure and is $\sim$ 4 times smaller than the value derived from our analysis. We also present the evolution of radio spectral indices with 1.06, 1.4, 4.86, 8.46 GHz flux density measurements. The spectral indices steepen to values of $-$1 around day 50 and continues till $\sim$ day 125, especially at higher frequencies (4.86/8.46), suggestive of electron cooling. During this period, the optical light curve is in the plateau phase. We estimate the cooling time scale for both IC and synchrotron cooling and interpret the steepening as a signature of IC cooling at the SN shock. SN\,2004dj is the only Type IIP SN with radio data covering two order of magnitudes in time and frequency and this allowed us to study this SN as a prototype of Type IIP SNe. We compare the properties of SN\,2004dj with other radio/X-ray bright Type IIP SNe and find that SN\,2004dj is a normal type IIP SNe with very similar CSM properties as that of SN\,1999em and SN\,2002hh. | 18 | 8 | 1808.00283 |
1808 | 1808.07837_arXiv.txt | We investigate the interactions of energetic hadronic particles (cosmic ray protons) with photons and baryons in protogalactic environments, where the target photons are supplied by the first generations of stars to form in the galaxy and the cosmological microwave background, while the target baryons are the interstellar and circumgalactic medium. We show that pair-production and photo-pion processes are the dominant interactions at particle energies above $10^{19}\;\! {\rm eV}$, while ${\rm pp}$-interaction pion-production dominates at the lower energies in line with expectations from, for example $\gamma$-ray observations of star-forming galaxies and dense regions of our own galaxy's interstellar medium. We calculate the path lengths for the interaction channels and determine the corresponding rates of energy deposition. We have found that protogalactic magnetic fields and their evolution can significantly affect the energy transport and energy deposition processes of cosmic rays. Within a Myr after the onset of star-formation the magnetic field in a protogalaxy could attain a strength sufficient to confine all but the highest energy particles within the galaxy. This enhances the cosmic ray driven self-heating of the protogalaxy to a rate of around $10^{-24}\;\! {\rm erg} \;\!{\rm cm}^{-3}\;\! {\rm s}^{-1}$ for a galaxy with strong star-forming activity that yields 1 core collapse SN event per year. This heating power exceeds even that due to radiative emission from the protogalaxy's stellar populations. However, in a short window before the protogalaxy is fully magnetised, energetic particles could stream across the galaxy freely, delivering energy into the circumgalactic and intergalactic medium. | \label{sec:introduction} High-energy hadronic particles heat the media along their propagation paths by collisional ionisation and hadronic interactions, with such action of cosmic rays (CRs) having been observed close-by in the Earth's atmosphere \citep[see e.g.][for reviews]{Ginzburg1996PhyU, Kotera2011ARA&A}. A substantial fraction of the energetic CRs detected on Earth are protons \citep{Abbasi2010PRL}, and their origins are likely to be extragalactic. The interactions of CRs in astrophysical environments have also been discussed in the literature in various contexts \citep[e.g.][]{Greisen1966PhRvL, Zatsepin1966JETPL, McCray1972ApJ, Nath1993MNRAS, Valdes2010MNRAS}. While research into CRs in solar-terrestrial settings has progressed a long way since their discovery, the role of ultra-high-energy (UHE) CR particles in galaxies and larger structures is less well understood -- even in systems where these energetic particles are evidently abundant, such as in star-forming galaxies \citep[see e.g.][]{Karlsson2008, Lacki2011ApJ, Lacki2012AIP, Wang2014} and galaxy clusters \citep[see e.g.][]{Brunetti2014}. It has been argued that CRs can regulate star-formation in a galaxy \citep[see][]{Pfrommer2007, Chen2016} and that their heating effect could drive large-scale galactic wind outflows \citep[see][]{Socrates2008, Weiner2009ApJ}. Not only do UHECRs alter the dynamical and thermal properties of their galaxy of origin, they can also transport energy from their host to the surrounding intergalactic medium. While CR heating and ionisation operate in nearby galaxies, conditions in the very distant Universe imply that CR processes are expected to be even more important when young galaxies were spawning their first generation of stars and luminous quasars first emerged. The effects of CRs in the high-redshift Universe are now beginning to be recognised \citep[see][]{Giammanco2005A&A, Stecker2006ApJ, Valdes2010MNRAS, Bartos2015, Sazonov2015MNRAS, Walker2015arXiv, Leite2017MNRAS}. CRs are products of violent astrophysical events, e.g.\ supernova (SN) explosions, gamma-ray bursts, large-scale shocks in colliding galaxies and galaxy clusters, and extreme environments in compact objects, e.g.\ fast spinning neutron stars and accreting black holes \citep[see][]{Berezinsky2006, Dar2008PhR, Kotera2011ARA&A}. As strong star-formation activity gives rise to frequent SN events, galaxies with active, ongoing star-formation are naturally strong CR sources. In these galaxies, the stars and their remnants supply the seed particles, and the shocks generated by the SN explosions and other violent events provide the needed mechanisms for accelerating the particles to very high energies, e.g. through Fermi acceleration processes \citep{Fermi1949PR}. Atomic matter can be ionised and excited by keV CR protons and this can heat the interstellar and intergalactic medium \citep[see e.g.][]{Nath1993MNRAS, Sazonov2015MNRAS}. However, compared to their higher-energy counterparts, a galactic population of sub-GeV CRs only contains sufficient energy to drive a relatively small heating effect, even if all of their energy were to be deposited within a reasonable timescale. This is because less than 1\% of the total CR energy is harboured in these lower energy particles \citep[see, e.g.][]{Benhanbiles-Mezhoud2013}. Higher-energy CRs (i.e. those with energy $\gtrsim 1 \;\!{\rm GeV}$) are much less engaged with atomic interactions. If ionisation and atomic excitation were the only means of energy exchange, a 0.5 GeV CR-proton would lose less than $2.5$\% of its initial energy on a Hubble timescale as it propagates across a cosmological baryonic density field, even when the density is maintained as that at redshift $z = 20$ \citep[see][]{Sazonov2015MNRAS}. However, the distance a CR particle above GeV energies can propagate in an interstellar or intergalactic medium is instead determined by hadronic processes. The Gresisen-Zatsepin-Kuzmin cutoff \citep{Greisen1966PhRvL, Zatsepin1966JETPL} is an example of such, where CRs are suppressed by their hadronic interactions with cosmic microwave background (CMB) photons. The impact of CRs depends on their interaction channels with the matter and radiation fields (microscopic particle physics) and also on the properties of the media through which the CRs propagate (microscopic astrophysics). How CRs affect the interior and exterior environments of their host galaxy is determined by the ability of the galaxy to act as a CR calorimeter (accounting for the galactic density profile and substructure), its regulation of CR diffusion, the properties of its global outflows and the entrapment of charged particles by its magnetic field \citep[see][]{Thompson2007, Tueros2014A&A, Kobzar2016ArXiv}. In this work we investigate the effects of energetic CR-particles in high redshift protogalactic environments. We adopt a phenomenological approach that sufficiently captures the essence of the relevant physics and astrophysics. Complexities such as the CR compositions at different stages of galactic evolution or the exact production and acceleration mechanisms of various types of primary particles in specific galactic components are finer details that will be left to future follow-up studies. Our focus is on the interactions of the particles with the radiation and baryon density fields and how efficiently energy is deposited along the particle propagations in the interstellar and intergalactic media. We organise the paper as follows. In \S~\ref{sec:interactions} we introduce the hadronic processes relevant to CR heating of the astrophysical media. In \S~\ref{sec:protogalaxies}, we specify the protogalactic model that provides the baryonic density profile, radiation properties and magnetic fields through which the propagations and interactions of the CRs are calculated. In \S~\ref{sec:discussion}, we show the results of our calculations, and we discuss their consequences and astrophysical implications in \S~\ref{sec:conclusions}. We particularly focus on those implications concerning propagation distances and energy deposition rates of the CRs and particle confinement in response to the evolution of the magnetic field of the host galaxy. For the remainder of this discussion, we refer to the UHECR protons as CRs, unless stated otherwise. For clarity we may also refer specifically to CR protons and CR electrons to differentiate between the two where appropriate, but when unspecified we refer to the proton component of cosmic rays. We do not consider primary CR electrons in detail as their energy loss rate is considerably more rapid compared to the CR protons. Therefore, their energy is deposited much more locally to their source and is less important in the global model we present here. | \label{sec:conclusions} \subsection{Astrophysical Implications} In this study we have demonstrated that CR heating power can exceed that due to conventional radiative heating provided by starlight and diffuse X-rays for conditions appropriate to protogalaxies exhibiting strong star-formation activity. This enhanced heating by CRs yields hotter proto-stellar gas, while CR scattering and interactions will induce additional turbulence throughout the ISM of the host \citep[see, e.g.][]{Niemiec2008ApJ, Bykov2011, Rogachevskii2012ApJ}. The resulting additional (thermal and turbulent) pressure support will lead to increased Jeans masses in star-forming regions \citep[see e.g.][]{Papadopoulos2011MNRAS, Papadopoulos2013}, with proto-stars pushed to higher initial masses \citep[see discussions in][]{Kuwabara2006ApJ, Ko2009ApJ, Hanawa2015ApJ, Hanawa2015PIAU, Kuwabara2015ApJ}. There are also suggestions that severe CR heating can quench star-formation entirely in host environments, perhaps even evaporating star-forming regions altogether in the most extreme cases \citep{Guo2008MNRAS, Gabor2011MNRAS}. Indeed, CRs may heat proto-stellar clumps throughout the host galaxy more effectively than radiation due to higher levels of energy deposition at the point of interaction and the possibility that they may be guided into star-forming regions by the threading magnetic fields \citep[see e.g.][]{Jokipii1966ApJ, Ferland2009}. In comparison with sub-keV X-rays and ultraviolet radiation, CRs are also less prone to shielding by the neutral material present in stellar nurseries and collapsing clouds and cores. Increased containment of CRs within the host galaxy would be expected to increase the rate of $\rm pp$-interactions driving the emission of $\gamma$-rays \citep{Abdo2010, YoastHull2015, Wang2014, Wang2018MNRAS} which is expected to be coupled with the CR heating effect (see section~\ref{sec:interactions} for details). This may also be the case for synchrotron emission following an enhanced number density of high energy CR secondary electrons injected into dense regions of the environment by the primary CR proton interactions. These secondaries may then interact with the strengthening galactic magnetic field \cite[see, e.g.][]{Schleicher2013A&A} and lose energy by synchrotron emission. At high redshift, the scattering of these energetic CR electrons off CMB photons is also likely to drive a level of X-ray emission. The resulting X-ray illumination in high redshift starbursts is first introduced in \citet{Schober2015MNRAS}, see also~\citet{Schleicher2013A&A}. On larger scales in the environment outside the host galaxy, the CR containment and the subsequent release of the CR energy may drive large-scale galactic outflows and fountains \citep{Shapiro1976ApJ, Bregman1980, Heesen2016MNRAS}. These outflows, in turn, may act as vehicles by which CRs are delivered into the circumgalactic medium and beyond into intergalactic space \citep{Heesen2016MNRAS}. This process may be a possible candidate for the pre-heating of the IGM in advance of cosmic reionisation \citep{Sazonov2015MNRAS, Leite2017MNRAS} at least in the proximity to a conglomeration of galaxies exhibiting violent star-forming activity. Additionally, the radiative emission from a host galaxy due to the contained CRs - in particular the X-ray emission resulting from inverse-Compton scattering of electrons~\citep{Schober2015MNRAS} - could begin to drive heating in the ISM of nearby neighbouring galaxies, possibly impacting on their own internal processes via X-ray heating. \subsection{Remarks} Observations, via Zeeman effect measurements \citep[see e.g.][]{Bergin2007, Koch2015, Ching2017ApJ} and dust emission polarimetry~\citep[e.g. see ][]{Planck2016bA&A} have revealed complex magnetic structures in the Galactic ISM. The alignment and the spatial structure of magnetic fields in and around dense and/or star-forming regions have shown great diversity and also vary over different length-scales in the same system \citep[see][]{Rao2009ApJ, Alves2011ApJ, Chen2012PIAU, Wang2015, Hull2017ApJ}. In the densest parts of some molecular clouds, where stars are spawned, the local magnetic field can reach levels of $\sim 1\;\!{\rm mG}$. \citep{Crutcher2010ApJ}. Such field strengths are strong enough to influence the propagation of CRs with energies even above a GeV, which will have gyration radii $\sim 10^{11}{\rm cm}$ (note that lower-energy MeV CRs will have smaller-scale gyration radii). The interaction of CRs with strong localised magnetic fields therefore needs to be modelled more thoroughly so as to determine their effects on CR scattering and propagation, and how the energy-selected CR containment by such localised magnetic fields can impact on both the local and global thermodynamics of the ISM of protogalactic environments. We have also examined the relative effect of starlight and X-ray heating. While intended as a comparison to the CR heating, there is an intriguing interplay between the different processes in addition to their relative importance in the thermodynamic context. For instance, the Compton scattering between the starlight photons with the high-energy electrons which are produced by the interactions of the CR protons may increase the level of X-rays inferred by stellar source counts, e.g. for SN remnants, X-ray binaries or colliding winds in stars. This effect may be particularly prevalent for violent star-forming protogalaxies at high redshifts and may considerably enhance the `indirect' CR heating impact on the interior and surrounding environment of the host. We have employed a simple parametric model as a demonstration of the importance of CR heating via hadronic processes in a protogalaxy. One of the crucial parameters in the model is the SN event rate, which determines the total power available for the CR heating process, be they confined or free-streaming. A uniform SN event rate is adopted in this study. The SN rate and magnetic fields are expected to co-evolve together if galactic scale magnetism in the protogalaxy is seeded by SN events. In this coevolution scenario, the containment of CRs in the initial and transient stage would be significantly different. However, one would not expect a substantial qualitative difference between including and excluding such effect once the magnetic field evolution has saturated. Although we have assumed a uniform SN event rate, the results obtained as such are nevertheless reasonable estimates. \subsection{Conclusions} In this study we have shown that the containment of CRs by a developing protogalactic magnetic field can arise within 10 Myr if employing a SN turbulent-dynamo driven model as outlined in \citet{Schober2013A&A}. Before this time, there is a window during which CRs may freely stream out of their host and into their surrounding environment. In the case of a protogalaxy actively forming stars at a rate of around $\sim 160~{\rm M}_{\odot} {\rm yr}^{-1}$ (leading to SN event rate of $\sim 1.0~{\rm yr}^{-1}$), the calculated containment amplifies the expected ISM heating effect due to the CRs from a level of around $10^{-29}$ ~erg ~cm$^{-3}$ s$^{-1}$ (which is lower than the conventional radiative heating mechanisms by stellar emission and X-rays) to around $10^{-24} ~{\rm erg~cm}^{-3}{\rm s}^{-1}$. Extrapolated from this level of CR heating within the central region of the protogalaxy, an extra-galactic heating rate of $\sim 10^{-37}~{\rm erg~cm}^{-3}{\rm s}^{-1}$ at around 10~kpc from the the protogalaxy, and $\sim 10^{-41}~{\rm erg~cm}^{-3}{\rm s}^{-1}$ at 100~kpc from the protogalaxy is also expected. The level of CR heating in the protogalactic environment obtained here also implies that CR heating in the IGM in the vicinity may also be non-negligible. This opens up many new questions about the impacts such an effect may have, and the degree to which it could be maintained. | 18 | 8 | 1808.07837 |
1808 | 1808.10036_arXiv.txt | We present a new X-ray study of NGC\,188, one of the oldest open clusters known in our Galaxy (7~Gyr). Our observation with the {\em Chandra X-ray Observatory} is aimed at uncovering the population of close interacting binaries in NGC\,188. We detect 84 sources down to a luminosity of $L_X$ $\approx$ 4$\times$10$^{29}$ erg~s$^{-1}$ (0.3--7~keV), of which 73 are within the half-mass radius $r_h$. Of the 60 sources inside $r_h$ with more than 5 counts, we estimate that $\sim$38 are background sources. We detected 55 new sources, and confirmed 29 sources previously detected by {\em ROSAT} and/or {\em XMM-Newton}. A total of 13 sources detected are cluster members, and 7 of these are new detections: four active binaries, two blue straggler stars (BSSs), and, surprisingly, an apparently single cluster member on the main sequence (CX\,33/WOCS\,5639). One of the BSSs detected (CX\,84/WOCS\,5379) is intriguing as its X-ray luminosity cannot be explained by its currently understood configuration as a BSS/white-dwarf binary in an eccentric orbit of $\sim$120 days. Its X-ray detection, combined with reports of short-period optical variability, suggests the presence of a close binary, which would make this BSS system a hierarchical multiple. We also classify one source as a new cataclysmic-variable candidate; it is identified with a known short-period optical variable, whose membership to NGC\,188 is unknown. We have compared the X-ray emissivity of NGC\,188 with those of other old Galactic open clusters. Our findings confirm the earlier result that old open clusters have higher X-ray emissivities than other old stellar populations. | \label{ch3_intro} X-ray production in single late-type stars is powered by a dynamo mechanism in their convective zones, which means that the faster the star rotates, the higher is its X-ray luminosity. According to the Skumanich law \citep{skumanich72}, the rotational velocities of single low-mass stars are proportional to the reciprocal of the square root of their age: as stars get older, they tend to slow down due to magnetic braking \citep{Pallavicini:1989p1112}, and correspondingly their X-ray emission decreases. Our Sun is one such old star ($\sim$4.5 Gyr) that has an X-ray luminosity of about $10^{26 - 27}$ erg s$^{-1}$ (0.1--2.4 keV; \citealt{Peres:2000p1107}). Such X-ray luminosities are nearly undetectable beyond distances of about a kiloparsec with the current generation of X-ray telescopes, not even with the sensitivity of the \textit{Chandra X-ray Observatory}. However, X-ray observations of old (age $\gtrsim$ 1 Gyr) open clusters have detected a rich population of X-ray sources that are associated with these clusters. Follow-up studies showed many of these sources to be close binaries of late-type stars that have been spun up by tidal interaction \citep{Belloni:1993p1075, Gondoin:2005p1051, Giardino:2008p1048, Gosnell:2012p685,vandenBerg:2004p1040, vandenBerg:2013p442, vatsvdb2017}. Such tidally interacting binaries, also known as active binaries (ABs), can have either two detached stars comprising the binary, or can have a contact or semi-detached configuration like W\,UMa and Algol binaries, respectively. X-ray sources in old clusters can also be accretion-powered, as in the case of cataclysmic variables (CVs), where a white-dwarf primary accretes matter from a late-type main-sequence donor. These CVs are typically found to the blue of the main sequence in the colour-magnitude diagram (CMD) due to the light from the accretion disk or stream, and possibly due to the contribution from the white dwarf itself. There are also some exotic X-ray sources found in old open clusters, like blue straggler stars (BSSs) and sub-subgiants (SSGs), however the origin of the X-ray emission from these sources, as well as the evolutionary status of these stars themselves, is not well understood (e.g.~\citealt{geller+17}). Old open clusters are very good laboratories to study binaries, such as ABs, CVs, SSGs and BSSs, as it is possible to obtain cluster membership information---and therefore age and distance---for a large number of sources with much less effort than for X-ray sources in the Galactic field. Also, the stellar densities of open clusters lie between those of dense globular clusters (GCs; $\gtrsim 10^4 M_{\odot}$ pc$^{-3}$) and the solar neighbourhood ($\sim0.1 M_{\odot}$ pc$^{-3}$). Hence, studying old open clusters aids us in understanding the role of stellar density in stellar and binary evolution. \citet{verbunt2000} found that the old open cluster M\,67 has a higher X-ray emissivity than most GCs. \citet{geea2015} demonstrated the elevated X-ray emissivity of two open clusters, viz. NGC\,6791 and M\,67, with respect to old stellar populations other than GCs, like dwarf galaxies and the local solar neighbourhood. In \citet{vatsvdb2017} we were further able to improve the statistics when we found that the old open cluster Collinder 261 (Cr\,261; age $\sim7$ Gyr) is also over-luminous in X-rays compared to a few dense GCs. To expand our understanding of open-cluster X-ray sources and the evolution of binaries in different environments, we are undertaking a survey with {\em Chandra} of old open clusters with ages between 3.5 Gyr and 10 Gyr. The observations are designed to reach a limiting luminosity of at least $L_{X} \approx 10^{30}$ erg s$^{-1}$ (0.3 -- 7~keV). In the current paper, we focus on the binary population of NGC\,188, which at an estimated age of $\sim$7 Gyr \citep{sarajedini99}, is one of the oldest open clusters in the Galaxy. It lies at a distance of 1650$\pm$50 pc, has half-mass and core radii of $r_h=8\farcm3\pm0\farcm6$ and $r_c=4\farcm4\pm0\farcm1$ respectively, and a reddening of $E(B-V) = 0.083$ \citep{Chumak:2010p1841}. NGC\,188 was first observed in X-rays using the $ROSAT$ Position Sensitive Proportional Counter (PSPC) with a flux detection limit of $\sim$10$^{-14}$ erg cm$^{-2}$ s$^{-1}$ in the energy range 0.1--2.4~keV \citep{Belloni:1998p1070}, equivalent to a luminosity limit of $\sim4\times10^{30}$ erg s$^{-1}$ for the quoted distance. The cluster was re-observed with {\em XMM-Newton} as a performance verification object with a luminosity threshold of $\sim$10$^{30}$ erg s$^{-1}$ (0.5--2.0~keV; \citealt{Gondoin:2005p1051}), but the cluster was not centred at the aimpoint during the observation, leading to asymmetric coverage of the cluster. Recent studies by the WIYN Open Cluster Survey (WOCS) have led to significant progress in our knowledge of NGC\,188. Membership of the cluster was established using proper-motion studies performed by \citet{platais2003} and stellar radial-velocity measurements performed by \citet{geller2008}. We use both of these studies for determining cluster membership of the likely and candidate optical counterparts to the {\em Chandra} sources detected in the study we present here. NGC\,188 hosts one of the best-studied populations of BSSs (e.g.\,\citealt{geller2011}, \citealt{gosnell2014,gosnell2015}). BSSs are stars that are bluer and brighter than the main-sequence turnoff point of a coeval population. They were first discovered in the globular cluster M\,3 by \citet{sandage53}, however, our understanding of their formation is still quite poor. There are currently three suggested scenarios for how BSSs are formed -- mass transfer in a binary system, merger of two (or more) stars due to a direct collision, and merger of the inner close binary induced by a tertiary companion in a hierarchical triple (see \citet{Davies2015book} for a review). Uncovering blue-straggler formation scenarios therefore contributes to our understanding of stellar encounters in clusters. Not all BSSs are expected to be sources of X-ray emission. However, if X-rays are observed from a particular BSS this could give away valuable clues to its current configuration (e.g.\,it may indicate the presence of a close binary in the system) and thereby constrain its past evolution. In Section \ref{ch3_obs_ana} we describe the observations and analysis. In Section \ref{ch3_results} we explain how we performed the source classification and we present our results. In Section \ref{ch3_discussion} we discuss our findings, and give details about two particularly interesting sources, including a BSS (WOCS\,5379) whose X-ray emission is not well understood. Section \ref{ch3_summary} is a summary of this work. \begin{figure} \includegraphics[clip=,width=1.0\columnwidth]{fig1.eps} \caption{$V$-band image of NGC\,188 and its surroundings from \citet{stetson2004}. The white dashed circle shows the central 20\arcmin\,of the $ROSAT$ PSPC observation of NGC\,188. The blue composite region shows the area covered by the {\em XMM-Newton} EPIC-PN and EPIC-MOS detectors. The black squares show the area covered by the {\em Chandra} ACIS-I chips. The red cross marks the centre of the cluster, and the outer and inner yellow circles mark the half-mass radius, $r_h$, and the core radius, $r_c$, respectively.} \label{ch3_fov} \end{figure} | \label{ch3_discussion} We conducted an X-ray study of the old open cluster NGC\,188 using {\em Chandra}. Our study is several times more sensitive than observations of the cluster done with {\em ROSAT} \citep{Belloni:1998p1070} and {\em XMM-Newton} \citep{Gondoin:2005p1051}. Unlike in these previous studies, the {\em Chandra} pointing was centred on the cluster centre, providing a more symmetric coverage of the cluster. Also, the positional accuracy of {\em Chandra} helps us in identifying candidate optical counterparts to the X-ray sources. Of the 84 X-ray sources detected by {\em Chandra,} 35 unique X-ray sources have a single candidate optical counterpart, while one X-ray source (CX\,76) has at least two candidate optical counterparts. Of the remaining 48 X-ray sources, 12 appear to have candidate optical counterparts in the $V$-band image by \citet{stetson2004}, however, these faint sources are neither included in the \citet{stetson2004} nor in the \citet{platais2003} catalogue. As shown by \citet{Mathieu:2004p602}, late-type main-sequence binaries in NGC\,188 show a transition around $P_{sp}=15.0$ days from circular orbits (below this period) to eccentric orbits (for longer periods). This demonstrates that at 7 Gyr (i.e.\,the age of NGC\,188) tidal interaction has had sufficient time to circularise orbits up to this cutoff period. Since tidal synchronisation of the stars' rotation occurs before circularisation (e.g.\,\citealt{Zahn:1989p220}), we expect late-type binaries with orbital periods at least up to 15 days to have achieved tidal coupling ($P_{rot} = P_{sp}$) and have enhanced levels of X-ray emission compared to single stars of similar age and mass. Indeed, all the ABs discussed in Section\,\ref{ch3_sec_ab} have periods shorter than 15 days, except CX\,13 for which the radial-velocity period is longer (discussed below). The X-ray activity of a late-type main-sequence star is related to its rotation period \citep{pizzolato03}. Generally, the X-ray luminosity increases with decreasing rotation period and follows the relation $L_X \propto P_{rot}^{-2}$. However, for short periods a star reaches a ``saturated'' regime, where $L_X$ no longer depends on $P_{rot}$, but instead flattens to $L_X/L_{bol} \sim 10^{-3}$ with $L_{bol}$ being the bolometric luminosity of the star \citep{randich98,Gudel:2004p12}. Stars with the shortest periods appear to emit X-rays at a level even lower than predicted by this saturated luminosity relation. Such stars are said to be ``super-saturated'' \citep{prosser96} and this behaviour has been ascribed to either a change in the dynamo action of the star or the reduction in the area of the active regions (see Sect.~5.4 in \citealt{Gudel:2004p12}). In Figure\,\ref{ch3_pvslx}, we plotted $L_X$ versus a measure of $P_{rot}$ for cluster members that were matched with a {\em Chandra} source to see if we observe this behaviour in NGC\,188. The periods used for the sources on the top panel were taken from the radial-velocity orbital solutions from \cite{geller2009} if available (CX\,13, CX\,26, CX\,53, CX\,84), and otherwise from variability surveys \citep{mochejska2008,zhang2004}. We can see that the overall behaviour of $L_X$ as function of period is in line with the trend described above, albeit with significant scatter. Moving from $\sim$7 days to shorter periods, $L_X$ reaches higher values until $\sim$1.3 days; this is similar to what we see for the ABs in M\,67 \citep{vandenBerg:2004p1040}. For shorter periods, $L_X$ starts diminishing again--this may correspond to the super-saturation regime. Three sources are outliers: CX\,13, CX\,26 and CX\,84. For CX\,13 and CX\,84 the period from radial velocities is different from the reported photometric period. If we plot the latter instead (bottom panel), these two sources are more in line with the trend described. CX\,13 and CX\,84 are described in more detail in Sections \ref{ch3_sec_cx13} and \ref{ch3_cx84}. The SSG CX\,26 is more X-ray luminous for its period compared to the overall trend; the spectroscopic period is the only measure for the rotation period available. There are nine cluster members within the {\em Chandra} FOV with orbital periods from \citet{geller2009} below $\sim15$ days, that were not detected in X-rays. Six of these sources could simply have X-ray luminosities below our detection limit as their periods are relatively long ($8$ days $< P_{sp} < 15$ days). In addition, WOCS\,5762 ($P_{sp}=6.5$ days) and WOCS\,5147 ($P_{sp}=6.7$ days) are found further away from the aimpoint where the X-ray detection limit is higher. However, we cannot explain the non-detection of WOCS\,5052. This spectroscopic binary lies closer to the aimpoint; it has a 3.85 day spectroscopic period and a near-circular orbit ($e=0.05\pm0.03$; \citealt{geller2009}). \subsection{CX\,13} \label{ch3_sec_cx13} The optical counterpart to CX\,13 (WOCS\,4705) is an eccentric ($e = 0.487\pm0.005$), double-lined spectroscopic binary with an orbital period of $P_{sp}\approx35.2$ days \cite{geller2009}. Its orbital parameters set this X-ray source apart from typical ABs, which have shorter orbital periods and circular orbits. Indeed, the location of this binary in the diagnostic diagrams of \citet[][Figures 5a and 5b]{verbuntphinney95} indicate the orbit is too wide for the primary to have circularised the orbit within the age of the cluster. However, since the time scale for tidal synchronisation is shorter than for circularisation, the rotation of one or both stars in this binary can already be faster compared to single (sub-)giants or turnoff stars in NGC\,188, which would explain the X-ray emission. The pseudo-synchronisation period, defined in \citet{Hut:1981p99} to be the spin period of the star if it corotates with the average orbital speed around periastron, is 13.1 days in this binary. This is a bit longer than the corotation period at periastron passage (10.6 days). CX\,13 may be an RS\,CVn-like system like the binary and X-ray source S\,1242 in M\,67, which is an eccentric ($e=0.66$) $P_{sp}=31.8$ day period binary on the subgiant branch. WOCS\,4705 is the optical variable V\,11. The observed 0.4 mag drop in brightness led \citet{kaluzny90} to postulate that V\,11 is a possible eclipsing RS\,CVn binary. Interestingly, \citet{geller2009}, based on their orbital ephemeris, found that this dip in brightness occurred at the orbital phase where an eclipse would indeed be expected to occur. Subsequent variability studies \citep{mazurkaluzny1990,zhang2004} only observed small-amplitude variations (0.03--0.08 mag). \citet{zhang2004} reported a possible photometric period of 1.2 days. We note that folding the published photometry for V\,11 on this period does not produce a smooth light curve; on the other hand, this period does give a better match to the $L_X$ versus $P_{rot}$ trend in Figure \ref{ch3_pvslx}. If real, the origin of this period is unexplained. CX\,13 lies within the core radius of the cluster, where the number of spurious matches between X-ray sources and photometric variables is expected to be $\sim0.013$, making the chance alignment of CX\,13 with a variable very unlikely. \subsection{CX\,84}\label{ch3_cx84} CX\,84 is a source located 3$\farcm$7 from the cluster centre that is identified with the BSS WOCS\,5379. From radial velocities \cite{geller2009} derived an eccentric ($e=0.24\pm0.03$) orbit with a period of $P_{sp}=120.21$ days. It is one of the seven BSSs in NGC\,188 for which FUV observations with the {\em Hubble Space Telescope} uncovered an excess flux compared to what is expected from the BSS alone, showing that the companion in these systems (all spectroscopic binaries) is a white dwarf \citep{gosnell2014,gosnell2015}. From the temperature of the white dwarf in WOCS\,5379 ($T\approx17,600$ K) it was derived that it formed only $\sim77\pm7$ Myr ago. \cite{gosnell2014} sketched a possible evolutionary scenario, in which the original secondary became a BSS after mass accretion from the white-dwarf progenitor. However, there are two aspects left unexplained by this evolutionary path. It fails to explain the highly eccentric orbit, as mass transfer via Roche-lobe overflow should circularise the orbit of the binary quite rapidly (although under certain circumstances this is not necessarily the case, see \cite{sepinsky2007,sepinsky2009}). Also, the nature of the binary components does not account for the X-rays: the white dwarf is not hot enough for thermal X-ray emission, and it is not obvious why the BSS would be an X-ray source. Another cluster BSS within the {\em Chandra} FOV, WOCS\,4540, with an age ($70\pm7$ Myr; \citealt{gosnell2014}) similar to WOCS\,5379, was not detected as an X-ray source. The key to solving this puzzle may lie in the optical variability that was detected for WOCS\,5379. \citet{kafhon2003} found this source to be a photometric variable (WV\,3) with a $V$ amplitude of 0.22 mag and a period of $P_{ph}=0.18148$ days (if the light curve has one maximum/minimum per period), or instead $P_{ph}=0.36296$ days if there are two maxima/minima per period. What is the origin of this variability? According to \citet{geller2009}, WOCS\,5379 lies outside the instability strip, ruling out pulsations as a likely explanation. Alternatively, the variability may be the signature of a short-period binary ``hiding'' in the system. The distribution of the data points around photometric minimum when folded on the shorter period \citep{kafhon2003} looks bimodal, which may be a sign that the actual period is $\sim$0.36 days and that the two minima are uneven. The presence of a short-period binary in the system would provide possible explanations for the X-rays (coronal activity, low-level mass transfer). On the other hand, it remains unclear what the exact configuration of such a hierarchical multiple system would be and whether any of the observed system components (the BSS and the white dwarf) are part of the short-period binary. \citet{geller2009} report that they were unable to derive a kinematic orbital solution for a period of 0.18148 days, and also do not find a sign of rapid rotation in the spectrum of WOCS\,5379. This implies that the BSS, which dominates the optical spectrum, is not part of the short-period binary. The white dwarf could have a close companion, perhaps similar to the pre-cataclysmic variables. But if mass transfer from the white-dwarf progenitor to the BSS progenitor in the wide orbit ended only $\sim77$ Myr ago (as inferred from the cooling age of the white dwarf), then this close companion entered the system only recently (in an encounter that perhaps also induced the eccentricity), complicating this speculative formation scenario even more. An extensive exploration of conceivable evolutionary scenarios is beyond the scope of this paper. In light of our findings, which suggest that CX\,84/WOCS\,5379 is not a binary but a triple or higher-order multiple, we conclude that its actual formation may be more complicated than the scenario of mass transfer between two binary components as sketched by \cite{gosnell2014}. In star clusters with a binary frequency exceeding 10\% (such as in NGC\,188 where the frequency of main-sequence binaries with periods up to $10^4$ days is 29\%$\pm$3\%; \citealt{geller2012}), binary--binary encounters typically dominate over binary--single-star encounters \citep{sigurdsson1993,leigh2011}, producing, among others, hierarchical triples. \cite{leigh2013} have demonstrated the significance of dynamical encounters involving triple systems, especially in open clusters. In fact, as has been argued by \cite{leigh2011}, a considerable fraction of the blue stragglers in NGC\,188 could have a dynamical origin. Such dynamical encounters can become quite complex: in low-mass star clusters such as old open clusters or low-mass globular clusters, the encounter durations are comparable to the time scale for another encounter to happen. As a result, the probability of an ongoing encounter being interrupted by a subsequent encounter can be up to a few tens of percent \citep{gellerleigh2015}. The same is found when binaries of hydrogen-rich, non-degenerate stars undergoing stable mass transfer are considered: in low-mass clusters, the encounter time scale is comparable to the duration of the mass-transfer phase, resulting in a significant probability that a mass-transfer binary could be disrupted or affected by other nearby stars \citep{leigh2016apj}. The outcomes of such events, or sequence of events, are difficult to predict for individual cases, but the point is that the possibility of dynamical interactions playing a role in the formation of CX\,84/WOCS\,5379 should not be neglected. We conclude that WOCS\,5379 is an interesting target for follow-up studies to improve the current light-curve and period measurements. Also, high--signal-to-noise high-resolution spectra would enable a search for activity or accretion signatures. By subtracting a high-quality template spectrum for the BSS, one could look for any additional components contributing to the optical spectrum. A similar technique was used to uncover a close binary in the spectrum of the BSS S\,1082 in M\,67 \citep{vandenBerg:2001p1001}, for which a long orbital period was found from radial velocities ($\sim1189$ days; \citealt{sandquist2003}) while at the same time eclipses on a 1-day period had been reported \citep{goranskij+92}. \subsection{Comparison With Other Old Star Clusters} We compare the number of X-ray sources that we find in NGC\,188 to those found in other old open clusters. To compare these results uniformly, we select sources brighter than $L_X\approx1\times10^{30}$ erg s$^{-1}$ (0.3--7~keV) that are found inside $r_h$. For NGC\,188, using a 2~keV MeKaL model to estimate the X-ray fluxes, we observe 55 such sources. Of these 55 X-ray sources, 9 are confirmed cluster members, and another 3 have uncertain membership classification (the remaining ones are likely AGNs). Among these, four are ABs (CX\,13, CX\,22, CX\,38, CX\,40), two are CV/AGN candidates (CX\,14, CX\,28), two are SSGs (CX\,26, CX\,49), one is a BSS (CX\,53), one an FK\,Com (CX\,4), and two sources have uncertain classification (CX\,33, CX\,48). We present these numbers in Table\,\ref{ch3_tab}, and compare them with four other old open clusters. We find that the addition of NGC\,188 to the sample confirms the trend that the number of X-ray sources appears to scale with mass in open clusters when the three clusters with the highest-quality membership information (M\,67, NGC\,188 and NGC\,6791) are considered. This is expected for a population that is (mostly) of primordial origin. The number of ABs does not show an obvious scaling with mass; the small source samples inhibit making any firm conclusions. It has been pointed out that globular clusters and other old stellar populations appear to be deficient in X-ray luminosity per unit mass, when compared to old open clusters (e.g.~\citealt{verbunt2000}, \citealt{vandenBerg:2013p442}, \citealt{geea2015}). X-ray emissivities of various old populations have been reported in the recent literature. Due to the variety of methods and energy bands adopted, it is not straightforward to compare the resulting values directly. The X-ray emissivities in \cite{geea2015} were computed in a consistent manner. To enable a comparison of our result for NGC\,188 with their set of X-ray emissivities of dwarf ellipticals and globular clusters, we have converted the X-ray emissivity of NGC\,188 to the 0.5--2 keV and 2--8 keV bands for the assumption of a 2~keV MeKaL model. We find that $L_X/M \approx 15.9 \times 10^{27}$ erg s$^{-1}$ $M_{\odot}^{-1}$ (0.5--2 keV), which is indeed higher than the X-ray emissivities in Table 2 from Ge et al. However in the 2--8 keV band, the extrapolated NGC\,188 emissivity is $L_X/M \approx 8.7 \times 10^{27}$ erg s$^{-1}$ $M_{\odot}^{-1}$, which is included in the range of X-ray emissivities in \cite{geea2015}. We caution that our adopted spectral model to compute X-ray fluxes is different than the one in their paper, which is a thermal model with a log-normal temperature distribution around a (fitted) peak temperature value. A uniform reanalysis of the X-ray data of old populations lies outside the scope of this paper, but we refer the reader to the forthcoming paper by Heinke et al.\,that makes an in-depth comparison of X-ray emissivities of globular clusters and those of other old stellar populations, including old open clusters. We illustrate the elevated X-ray emission of old open clusters with concrete numbers. NGC\,188 has a mass of $2600\pm460$ $M_{\odot}$ \citep{geller2008} and has two members (the FK\,Com star CX\,4, and the subgiant AB CX\,13) with $L_X\gtrsim5\times 10^{30}$ erg s$^{-1}$ (0.3--7 keV) within $r_h$. In M\,67, with a mass of 2100$^{+610}_{-550}$ \citep{geller2015} there are three members above this luminosity limit: a SSG, a BSS and an AB. Finally, in NGC\,6791 (5000--7000 $M_{\odot}$; \citealt{platea11}), seven to eight members are detected above this luminosity limit; they are a mix of CVs, SSGs, ABs and an apparently single red giant. In contrast, the two sparse globular clusters Terzan\,3 and NGC\,6535 with masses $\sim$12\,000 $M_{\odot}$ were also observed down to $\sim5\times 10^{30}$ erg s$^{-1}$ (0.3--7 keV) with {\em Chandra}, and zero sources were detected inside $r_h$ (A.\,Kong, private communication). Several factors can contribute to the enhanced X-ray emissivity of old open clusters, including the faster rate at which stars are lost from open clusters compared to globular clusters as a result of their shorter relaxation times (reducing their mass more severely), the higher rate of dynamical destruction in globular clusters of certain types of binaries that contribute to the X-ray source populations of open clusters, differences in age, and differences in metallicity (see also the discussion in Section 5.2 of \citealt{vatsvdb2017}). Further studies of more clusters with a range of properties are needed to gain more insight. We present the results of a new X-ray study of the 7-Gyr old open cluster NGC\,188. {\em Chandra} detected 84 sources above $L_X\approx4\times10^{29}$ erg s$^{-1}$ (0.3--7 keV), of which 73 lie within the half-mass radius. Of the thirteen cluster members that were detected by {\em Chandra}, seven are new X-ray detections. The X-ray source population of this cluster is a mix of ABs, SSGs, BSSs, and a single FK\,Com star. While some of these X-ray source types are frequently seen in other old open clusters, we also found a few surprises. CX\,33 is an apparently single cluster star on the main sequence, and we do not understand its X-ray emission. CX\,13 (V\,11) is an AB that may be on its way to being circularised. The orbital parameters of the BSS and white-dwarf binary WOCS\,5379 already were puzzling given the combination of an eccentric orbit in a post--mass-transfer system, but the detection of X-rays (CX\,84) makes this system even more intriguing. It raises the question whether the short-period variability that has been reported for this star is coming from a close binary inside the system, which would also explain the X-ray emission. Considering the overall X-ray population in NGC\,188, we confirm our findings that the X-ray emissivity of old open clusters is elevated compared to other old stellar populations. | 18 | 8 | 1808.10036 |
1808 | 1808.05235_arXiv.txt | We analyze 168 \emph{Swift} monitoring observations of the nearest broad absorption line quasar Mrk\,231 in the UV and X-ray bands, where we detect significant variability in the UV ($\sim$2246\AA) light curve with a null probability of $4.3\times10^{-10}$ for a constant model. Separately, from an archival sample of \emph{Swift} observed active galactic nuclei (AGN), we measure the relation between UV excess variance and luminosity, finding that the normalized UV excess variance decreases with luminosity. Comparing to this mean relation, the normalized UV excess variance of Mrk\,231 is smaller, however within the scatter characterising the full population. The upper limit of the X-ray excess variance is consistent with other AGN. The power spectrum density of the UV light curve can be well fit by a power law model with a slope of $1.82\pm0.14$ between $10^{-7.5}$ and $10^{-6}$~Hz, consistent with those for typical AGN, with no obvious quasi-periodical oscillation peaks. The UV variability and its power spectrum suggest that a significant amount of the UV emission of Mrk\,231 is from the accretion disk. The consistencies in the normalized UV variability and the shape of the power spectrum density between Mrk\,231 and other normal AGN suggest that the origin of UV variability of broad absorption line quasars is similar to other AGN, and dust scattering at large scales such as the torus is not a dominating process for the UV emission of Mrk\,231. Significant scattering, if present, is constrained to smaller than $\sim$10 light days. We perform lagged correlation analysis between the UV and X-ray light curves and find the correlation insignificant within the present data. | \label{sec:1} Timing analysis is a powerful tool to study the properties of active galactic nuclei (AGN), since they in general exhibit variability on a vast range of time scales and wavelengths. In situations when the data is sparse, the noise removed variance of the light curve is often used as a first step to characterize the variability \citep[e.g.,][]{Nandra1997, Vaughan2003}. In general, it is found that luminous AGN have relatively smaller variability amplitudes compared to the less luminous ones \citep[e.g.,][]{Nandra1997,1998ApJ...501L..37P,2000MNRAS.315..325A,2002MNRAS.330..390M,2004ApJ...605...45D}. When the data quality is high and densely sampled, Fourier analysis can be applied to obtain the power spectrum density (PSD) of the light curve. Previous studies found that the PSD of AGN is generally composed of a single or multiple segments of power laws \citep[e.g.,][]{2002MNRAS.332..231U,2003ApJ...593...96M,2011ApJ...743L..12M,2013MNRAS.430L..49M,2014ApJ...795....2E}. There are also situations in between, and the data is used to test a particular PSD model, e.g., the damped random walk model, which has a PSD power law slope of $-2$ but flattening at low frequencies. Many studies have concluded that the optical AGN variability is consistent with the damped random walk model \citep[e.g.,][]{2009ApJ...698..895K, 2010ApJ...721.1014M, 2010ApJ...708..927K}. Most previous PSD analyses of AGN focus on the visible or X-ray light curves, and in very few cases densely-sampled UV light curves are available for the PSD analysis. This is especially true for the sample of broad absorption quasars, which show broad, blue-shifted, absorption troughs in their spectra, and it is estimated that broad absorption line quasars contribute to 20--40\% of Type I quasar population intrinsically \citep[e.g.,][]{2008ApJ...672..108D, 2008ApJ...687..859S,2011MNRAS.410..860A}. Broad absorption line quasars are also observed to be either X-ray weak or intrinsically X-ray weak \citep[e.g.,][]{1996ApJ...462..637G,2000ApJ...533L..79M,2002ApJ...567...37G,2003AJ....126.1159G, 2011ApJ...737...46M, 2014ApJ...786...58M}. Correlations between multi-band continuum light curves can reveal the origin of the disk variability. As reviewed by, e.g., \citet{MH16}, the origin of the UV-optical variability in AGN is still a matter of debate. Leading theories broadly suggest either a reprocessing of X-ray emission by the accretion disc, or simply intrinsic variability of the thermal emission by the disc itself. Reverberation mapping has been adopted as a valid discriminator between these two proposals. For example, evidence of a time lag between the UV/optical variations occurring later than the X-ray ones would be in support of the former scenario. The ultra-luminous infrared galaxy Mrk\,231 hosts the nearest quasar at a redshift of $z=0.042$, where the AGN contribution of the bolometric luminosity is estimated to be $10^{46}$~\lumin\,\citep[e.g.,][]{2009ApJS..182..628V}. Mrk\,231 is also a broad absorption line quasar with low-ionization lines including Fe absorption troughs \citep{1972ApJ...173L.109A}, and thus is classified as an FeLoBAL. FeLoBALs contribute to 2\% of Type I quasar population intrinsically \citep{2012ApJ...757..180D}. Mrk\,231 is intrinsically X-ray weak as revealed by \emph{Nustar} observations, the X-ray absorption column density is constrained to be heavily absorbed but Compton-thin ($\nh = [1.2\pm0.3] \times 10^{23}~\cmsq$), and the intrinsic X-ray spectrum is flat with a power-law photon index of $\Gamma=1.4$ \citep{2014ApJ...785...19T}. Overall, Mrk\,231 exhibits many unusual features among quasars. Recently, an additional unusual spectral feature was discovered in the ultraviolet (UV) bands \citep{2013ApJ...764...15V,2014ApJ...788..123L,2015ApJ...809..117Y}. The continuum of Mrk\,231 drops dramatically from the optical to the near ultraviolet (NUV) band; however, the drop stops at the far ultraviolet (FUV) band, and FUV emission is detected all the way to $\sim$1000\AA. There are several models to interpret the anomalous UV spectrum of \mrk. It is potentially possible to distinguish these models from the UV variability of \mrk, in addition to measure the overall UV variability characteristics of a broad absorption line quasar. \citet{2015ApJ...809..117Y} measured UV emission of Mrk\,231 by analyzing the archival International Ultraviolet Explorer (\emph{IUE}) and Hubble Space Telescope (\emph{HST}) data, and found it to be variable. \citet{2016ApJ...829....4L} and \citet{2016ApJ...825...42V}, however, argued that the UV emission of Mrk\,231 may not vary based on analyses on \emph{HST} observations only. It is therefore important to check whether the UV emission from Mrk\,231 varies or not. Most of the limited UV observations of Mrk 231 are spectroscopic. These include three spectroscopic measurements by \emph{IUE} between 1978--1980 and a few spectroscopic measurements by \emph{HST} from 1996 to 2014. Recently, the Neil Gehrels Swift Observatory (hereafter \swift) has made a large number of imaging UV observations on Mrk\,231 during the period from 2012 to 2016, providing accurate photometric flux measurements. In this paper, we analyze the \emph{Swift} UVOT and X-ray data to study the variability of Mrk\,231 in these bands. In Section\,\ref{sec:2}, we describe the method to process \emph{Swift} data. We analyze the light curves obtained from \emph{Swift} UV and X-ray data in Section\,\ref{sec:3}. Conclusions and discussion are given in Section\,\ref{sec:4}. | \label{sec:4} In this paper, we have analyzed the UV and X-ray observations of Mrk\,231 core by \emph{Swift} and found that its UV emission varies significantly based on the $\chi^2$ test result. The detection of UV variability for Mrk\,231 is consistent with our previous analysis in the FUV band ($\sim$1030\AA) based on the archival \emph{IUE} and \emph{HST} data \citep{2015ApJ...809..117Y}, and in contrast with other analyses claiming that the UV flux of Mrk\,231 is a constant. UV variability from the accretion disk is a characteristic of many AGN observed by \emph{Swift} \citep[e.g.,][]{2010ApJS..187...64G, 2012ApJS..201...10W, 2013A&A...550A..71V}. The PSD of the UV light curve of Mrk\,231 is well fit by a power law model, another characteristic of AGN variability, where the slope is measured to be $1.8\pm0.3$ for frequencies below the flattening caused by white noise. Very few AGN have the power spectrum measured in the UV band, and thus we compare with those measured in the optical and X-ray bands. In the frequency range between $10^{-5.5}$ to $10^{-7.5}$\,Hz, the typical slope for the PSD of AGN is $\beta \sim 1-2$ in the X-ray band \citep[e.g.,][]{2002MNRAS.332..231U, 2003ApJ...593...96M, 2013MNRAS.430L..49M} and $\beta \sim 2-3$ in the optical band \citep[e.g.,][]{2009ApJ...698..895K, 2011ApJ...743L..12M, 2014ApJ...795....2E}. The UV PSD slope measured in Mrk\,231 is consistent with the damped random walk model of AGN optical variability \citep[e.g.,][]{2009ApJ...698..895K, 2010ApJ...721.1014M, 2010ApJ...708..927K}. Since Mrk\,231 is a broad absorption line quasar, an FeLoBAL in particular, dust is expected to be present in the central engine. However, the amount of dust, its location, and how it affects the UV emission are all uncertain. If dust scattering is important for the UV emission of \mrk, we expect that the PSD will be significantly suppressed at frequencies higher than that corresponding to the dust scattering scale. For broad line regions, the size scale is $\sim$60 light days for $10^8\msun$ black holes \citep[e.g.,][]{2011A&A...528A.100S}, and for the dusty torus, typical estimates are $\sim$300 light days \citep[e.g.,][]{2006ApJ...639...46S} for quasars. These scales correspond to light-crossing times with frequencies between $4\times10^{-8}$ and $2\times10^{-7}$~Hz, and we do not measure significant suppression of variability power above these frequencies. Instead, the PSD extends to $10^{-6}$~Hz following the typical power-law model for AGN. This suggests that dust scattering at large scales is not a dominating process for the UV emission of \mrk. The high frequency PSD measurements allow us to constrain the UV emission or scattering scales smaller than $\sim$10 light days. The upper limit of the X-ray variability measured for \mrk, with a massive black hole of $\sim 1.5\times 10^8 \msun$ based either on our $H_\alpha$ line width measurement from the optical spectrum reported in \cite{2014ApJ...788..123L} and the virial mass estimator of \cite{2005ApJ630122G} or the BBH fit of \cite{2015ApJ...809..117Y}, is consistent with the general finding that (short) time-scale variability is inversely proportional to black hole mass \citep[][and references therein]{2001MNRAS.324..653L,2013MNRAS.430L..49M}. Our findings pointing to an anti-correlation between amplitude of UV variability and optical luminosity is also consistent with the idea that variability is anti-correlated to mass accretion rate, though most probably additional physical drivers can complicate this basic picture \citep[e.g.,][and references therein]{Ive14}. The time lag between X-ray and UV-optical fluxes, with the latter usually lagging behind, is broadly consistent with the idea of reprocessing of higher energies photons by an accretion disk. At longer wavelengths, the lag $\tau$ has been shown to increase steadily with increasing wavelength as $\tau\propto \lambda^{\beta}$, with $\beta \gtrsim 1$, as expected from an extended and/or non-uniform disc \citep[][and references therein]{MH17}. Extrapolations of lags to shorter wavelengths/higher energies, does not always agree with the measured X-ray to UV lag \citep{MH17}. \citet{MH14} find long-term UV/optical variations are not necessarily paralleled in the X-rays, suggesting an additional component to the UV/optical variability possibly arising from accretion rate perturbations. \citet{Pal18} studied the variability in the Seyfert 1 galaxy NGC 4593, finding evidence for both a highly variable component such as hard X-ray emission, and a slowly varying disc-like component. Their data suggest the observed variation in longer wavelengths to be due to X-ray reprocessing. \citet{Cac18} monitored the Seyfert 1 galaxy NGC 4593 broadly confirming a lag spectrum $\tau\propto \lambda^{4/3}$ relation consistent with the standard thin disk model. However, they also suggest that larger disk sizes and emission from the Broad Line Region should also be considered as essential components to properly model AGN lag spectra. We do not detect significant lag between the UV and X-ray light curves from \emph{Swift} observations of Mrk\,231 mostly because of the low signal-to-noise ratio of the X-ray light curve. A similar monitoring campaign with XMM-Newton will be much more promising to detect the UV and X-ray lag in \mrk, since the UVOT instruments onboard of \emph{Swift} and XMM-Newton are similar while the throughput of X-ray mirror of XMM-Newton is an order of magnitude higher. We measured the normalized excess variance of Mrk\,231 in the UV band and compared with the values from a sample of nearby AGN observed by \swift. We performed a linear fit to the normalized excess variance measured in the UV band with the 5100\AA\ luminosity including intrinsic scatter, and we found that Mrk\,231 is below the mean relation from $0.3$ to $0.6$\,dex, but still consistent with other AGN considering the intrinsic scatter of the sample ($\sim 0.6$\,dex). Although within the scatter, we discuss several other factors that could contribute to \mrk's relatively low variability amplitude compared to the mean relation. First, Mrk\,231 is a broad absorption line quasar, and the presence of the disk wind may suppress the observed optical/UV variability from the disk. \citet{Conno14} suggest that the complex X-ray variability observed in the Seyfert 1.8 galaxy NGC 1365, could also be interpreted with a wind model in which the launch radius moves out with increasing X-ray luminosity. In addition, it is possible that Mrk\,231 is accreting in the super-soft state as proposed by \citet{2014ApJ...785...19T}, which predicts a low variability amplitude and a very hard X-ray spectrum as well. Indeed, \citet{Conno16} from the X-ray spectral variability of 24 local AGN from the Palomar sample of nearby galaxies, found that AGN with low accretion rates show hardening with increasing count rate, converse to the softer-when-brighter behaviour normally observed in AGN with higher accretion rates. Finally, the contribution of large-scale UV emission can be significant and dilute the variability from the central engine, as indicated by the FUV \hst\ image that Mrk\,231 is slightly extended \citep{2016ApJ...829....4L}. If the large-scale UV emission contributes to 30-50\% of the observed UV luminosity, the normalized excess variance of Mrk\,231 will be on the mean relation. As recently reviewed by \citet{P17}, high column densities of circumnuclear absorbing material can also significantly affect the fraction and extension of the X-ray variability in AGN. \citet{Gonza18} found non-trivial variations in 19 out of 22 X-ray-selected AGN, concluding that obscuration along the line of sight is an important parameter in shaping the observed correlations between, e.g., black hole mass, accretion rate, and break frequencies. \citet{H15} find from a sample of 26 Seyfert 2 galaxies that short-term X-ray variability is mostly associated to Compton-thin sources, which should more safely arise from variations in the nuclear source. In their sample UV variability on longer timescales seems not to be affected by the level of line-of-sight obscuration. \citet{Sanchez17} presented variability for a large sample of X-ray-selected AGN in COSMOS with different levels of obscuration. They found that broad-line AGN have a larger fraction of variable sources than narrow-line ones, and that X-ray-classified unobscured AGN tend to have a lower fraction of variable sources with respect to optically-classified unobscured AGN, possibly due to differences in the origin of the obscuration. Combining the variability, PSD, and excess variance analysis, we conclude that a significant fraction of the observed UV emission of Mrk\,231 in uvm2 band centered at 2246\AA\ is from the accretion disk. These dense \emph{Swift} UV monitoring data of Mrk\,231 provide the first opportunity to characterize the UV variability from a broad absorption line quasar in detail, although Mrk\,231 belongs to the rarer sample with Fe absorption troughs. We found that its fractional variability and shape of the PSD are consistent with other normal AGN. This suggests that the origin of UV variability of broad absorption line quasars is similar to other AGN. The central wavelength is shorter than the sharp drop-off observed in the \hst-COS spectrum around 3000\AA. We discuss the implications for the models proposed to explain the sharp drop-off in the UV continuum of \mrk. For the disk leakage model \citep{2013ApJ...764...15V, 2016ApJ...825...42V}, the sharp drop-off at the optical-to-NUV bands is due to an increase of extinction at shorter wavelengths by the intervening absorbers that cover the optical-to-UV continuum emission region. The flat continuum spectrum at FUV is due to the emission leaked out of holes in the spherically distributed absorbers ($\sim$5\% of the sky area). The leaked emission should still carry the variability signal from the accretion disk, in addition to the scattering dust. The UV PSD measurement from this paper has constrained the scattering to be smaller than $\sim$10 light days. The simple BBH model of \citet{2015ApJ...809..117Y} predicts UV variability broadly consistent with the observational results presented in this paper, since the UV emission is from an accretion disk in this model. However, the predicted amount of ionization photons differ significantly depending on detailed model assumptions \citep{2015ApJ...809..117Y, 2016ApJ...829....4L}. In addition, the BBH model may suggest a quasi-periodical variation of the UV emission around $\sim$1.2\,yr for \mrk; however, this frequency is outside of the range that can be probed by the current data set. The PSD of the UV light curve of Mrk\,231 is consistent with a single power-law with no obvious peaks from quasi-periodical variations. \citet{2014ApJ...788..123L} proposed that the sharp drop-off in the continuum is caused by a special dust extinction curve, which will completely block the FUV and a portion of UV emission from the accretion disk, and thus the model places the UV emission at large scales, such as the starburst contribution \citep{2014ApJ...788..123L}, or FUV photons scattered out of larger scale media \citep{2016ApJ...829....4L}. The \emph{Swift} detection of UV variability and the UV PSD shape show that the UV continuum is consistent with emission from the accretion disk and thus incompatible with the special dust extinction model that completely blocks the FUV/UV emission from the disk. The nature of UV emission of Mrk\,231 is still a mystery, and more observations are needed to better understand this prototype broad absorption line quasar. For example, in light of the analysis of this paper, measuring the PSD to even higher frequencies will put stronger constraints on the size and origin of the UV emission of \mrk. We acknowledge the financial support from the NSF grant AST-1413056. This work is also partially supported by the National Key Program for Science and Technology Research and Development (Grant No.\ 2016YFA0400704), the Strategic Priority Program of the Chinese Academy of Sciences (Grant No.\ XDB 23040100), and the National Natural Science Foundation of China under grant Nos.\ 11690024 and 11390372. FS thanks Ian McHardy for useful discussion. We thank Kayhan Gultekin for helpful discussion. \begin{table*} \centering \caption{Counts rates of Mrk\,231 core in the X-ray (0.3--8~keV) band from \emph{Swift} XRT observations and AB magnitude of Mrk\,231 core in uvm2 filter form \emph{Swift} UVOT observations. Table~\ref{tab:obs} is published in its entirety in the electronic edition of the journal. A portion is shown here for guidance regarding its form and content. \label{tab:obs}} \begin{tabular}{cccccc} \hline\hline Obsid & Obs.\ Start Time & XRT & X-ray (0.3--8\,keV) & UVM2 & UVM2 \\ & JD$-$2457414.53310 & Exposure (s) & $10^{-3}$ (counts/s) & Exposure (s) & AB Mag \\ \hline 00032530003 & 0 & 642.23 & $ 10.53 \pm 5.04 $ & 712.16 & $ 17.33 \pm 0.03 $ \\ 00032530004 & 2.0560 & 807.67 & $ 8.52 \pm 3.97 $ & 974.68 & $ 17.32 \pm 0.03 $ \\ 00032530005 & 4.5230 & 170.18 & $ -0.37 \pm 5.90 $ & 181.62 & $ 17.33 \pm 0.06 $ \\ 00032530006 & 6.0496 & 810.47 & $ 13.38 \pm 4.78 $ & 992.62 & ... \\ 00032530007 & 8.7092 & 889.76 & $ 14.54 \pm 4.66 $ & 945.25 & $ 17.35 \pm 0.03 $ \\ 00032530008 & 10.5587 & 789.24 & $ 6.34 \pm 3.52 $ & 1045.85 & $ 17.32 \pm 0.03 $ \\ 00032530010 & 14.6152 & 901.12 & $ 9.92 \pm 3.94 $ & 966.21 & $ 17.31 \pm 0.03 $ \\ 00032530011 & 16.6239 & 639.18 & $ 5.99 \pm 4.06 $ & 789.94 & $ 17.31 \pm 0.03 $ \\ 00032530012 & 18.8773 & 483.62 & $ 8.15 \pm 5.31 $ & 584.97 & ... \\ 00032530013 & 20.8732 & 924.37 & $ 19.17 \pm 5.20 $ & 990.01 & $ 17.32 \pm 0.03 $ \\ 00032530014 & 22.6711 & 877.29 & $ 14.66 \pm 4.74 $ & 931.68 & $ 17.31 \pm 0.03 $ \\ \hline \end{tabular} \end{table*} | 18 | 8 | 1808.05235 |
1808 | 1808.00404_arXiv.txt | We present a 3D study of the formation of {refractory-rich} exospheres around the rocky planets HD219134b and c. These exospheres are formed by surface particles that have been sputtered by the wind of the host star. The stellar wind properties are derived from magnetohydrodynamic simulations, which are driven by observationally-derived stellar magnetic field maps, and constrained by Ly-$\alpha$ observations of wind mass-loss rates, making this one of the most well constrained model of winds of low-mass stars. The proximity of the planets to their host star implies {a high flux of incident stellar wind particles, thus the sputtering process is sufficiently effective} to build up relatively dense, {refractory-rich} exospheres. The sputtering releases refractory elements from the entire dayside surfaces of the planets, with elements such as O and Mg creating an extended neutral exosphere with densities larger than 10~cm$^{-3}$, extending to several planetary radii. For planet `b', the column density of O{\sc i} along the line of sight reaches $10^{13}$~cm$^{-2}$, with the highest values found ahead of its orbital motion. This asymmetry would create asymmetric transit profiles. To assess its observability, we use a ray tracing technique to compute the expected transit depth of the O{\sc i} exosphere of planet `b'. We find that the transit depth in the O{\sc i} $1302.2$~\AA\ line is $0.042\%$, which is a small increase relative to the continuum transit ($0.036\%$). This implies that the sputtered exosphere of HD219134b is unlikely to be detectable with our current UV instruments. | Winds {from low mass, main-sequence stars} are formed from streams of {charged} particles that outflow from stars and, thus, permeate the interplanetary medium. As they make their way towards the interstellar medium, stellar wind particles drag along the stellar magnetic field. This magnetised plasma then interacts with orbiting exoplanets in a similar way as the solar wind interacts with solar system planets. The nature of the wind-planet interaction mainly depends on whether a planet is magnetised or not and whether it has a thick gaseous atmosphere or not \citep[e.g.][]{russell2016}. The solar system offers us some illustrations of different types of interactions. For example, the Earth's atmosphere is shielded from the direct interaction with the solar wind due to the presence of an extended (10--15\,$R_{\oplus}$) magnetosphere, which carves a cavity in the solar wind plasma, deflecting it around our magnetosphere. The magneotsphere of Mercury, which has a much weaker magnetic field, also carves a cavity in the solar wind, however with a significantly smaller size of only about 1.5 Mercury radii \citep{2013pss3.book..251B}. Mars, instead, is an example of a planet whose atmosphere directly interacts with the solar wind, creating an induced magnetosphere \citep{2011SSRv..162..113B}. In analogy to the interactions between the solar wind and the solar system planets, exoplanets {will also} interact with the winds of their host stars. Important differences can however exist, as both the architecture of known exoplanetary systems and the properties of the host stars can be significantly different from those of the solar system \citep{2015MNRAS.449.4117V}. For example, close-in exoplanets orbit at very short distances from their hosts, where the stellar wind is denser and the embedded magnetic field is stronger, when compared to planets orbiting at large distances. For this reason, close-in planets usually interact with harsher stellar wind environments than farther out planets. Likewise, planet-hosting stars might be quite different from the Sun (e.g., more magnetically active), such that even planets that are not necessarily too close from their host stars might interact with winds of significantly different properties \citep[e.g.,][]{2013A&A...557A..67V}. For planets similar to Mars (i.e., non-magnetised and with a thin, yet collisional atmosphere) or Mercury (i.e., weakly-magnetised and with a tiny, non-collisional atmosphere), the stellar wind interacts directly with either their atmosphere or solid surface. Although lacking a substantial atmosphere, bodies like Mercury may hold a tenuous (i.e., non-collisional) gaseous envelope, forming their exospheres. This exosphere is made up of particles sputtered from the surface by precipitating solar wind protons and following ballistic orbits around the planet. Photoionisation of these neutral particles creates an ion population in addition to the ions directly ejected from the surface. The continuous supply of the exosphere is sustained by various surface-release processes, like photon stimulated desorption, thermal evaporation, micrometeoroid impact, and ion sputtering. Among these processes, sputtering is considered to be the most energetic mechanism, leading to particles with energies of up to several hundreds of eV, distinctly exceeding the escape energies of many species at the surface of Earth-like planets. For Mercury, the surface, exosphere, and magnetosphere, together with the solar wind, constitute a complex and strongly coupled system dominated by the interaction of the neutral and ionised particles with the surface and the magnetospheric plasma. The upcoming ESA mission {\it BepiColombo} to Mercury is specifically devoted to the investigation of this highly dynamic and complex hermean environment \citep{2005SSRv..117..397M, 2007SSRv..132..433K}. In the case of close-in airless exoplanets, sputtering of their surfaces may be stronger than at Mercury, thus raising the question whether their expected exospheres might be observable. We study here the multi-planetary system \hd, which hosts six detected planets to date. Five of them orbit the star at separations smaller than $0.4$~au, while one distant gaseous giant planet orbits at $\approx 3$~au. The host star is a K3 dwarf, with an estimated age of $11.0\pm2.2$~Gyr \citep{2017NatAs...1E..56G} and an average large-scale surface magnetic field of $\approx 2.5$~G (\citealt{paper1}, henceforth \citetalias{paper1}). The two inner-most planets, \hd\,b and c, which are likely tidally locked, are {observed in transit} and present Earth-like densities \citep{2017NatAs...1E..56G}. {Their low-gravities suggest that both planets lost through escape their primary hydrogen-dominated atmospheres, presumably accreted during formation. The removal of their primary atmospheres is likely to have happened while the system was still young and the star was active \citep[e.g.,][]{2011A&A...532A...6S,2015A&A...577L...3T}.} Following the escape of their primary atmospheres, most likely the solidification of the magma oceans led to the formation of steam CO$_2$-dominated atmospheres \citep{2008ApJ...685.1237E,2012AREPS..40..113E,2018A&ARv..26....2L}. Depending on the past evolution of the high-energy {(X-ray and extreme ultraviolet, collectively called XUV henceforth)} stellar radiation, the planets may have lost also the CO$_2$-dominated secondary atmospheres \citep{2009ApJ...703..905T}, leaving behind the bare planetary surfaces directly interacting with the stellar wind. In this work, we start from the assumption that both planets have lost their CO$_2$-dominated atmosphere and do not host a significant magnetic field. Under these assumptions, the close proximity of both planets to the host star, and thus high {proton flux} of the incident wind, make the surface sputter, similarly to what occurs on Mercury. Some of the sputtered planetary particles would then ionise, forming mostly neutral and ionised Na, O, Si, and Fe atoms \citep[e.g.,][]{2009ApJ...703L.113S,2011ApJ...742L..19M,2015P&SS..115...90P,2016ApJ...828...80K}. The structure and velocity of the material escaping from the planet would then be controlled by the stellar wind properties, radiation pressure, and interplanetary magnetic field carried by the stellar wind. Here, we use state-of-the-art models of stellar wind and wind-induced sputtering to investigate the effects that the wind of \hd\ has on building up exospheres on the two inner-most planets, how the wind interacts with it and whether these atmospheres can be observed. This paper is organised as follows. In Section~\ref{sec.wind}, we present our stellar wind model, which uses surface magnetic field maps derived from the Zeeman-Doppler Imaging technique \citepalias{paper1}. Our wind model is constrained by the observed mass-loss rate derived from Ly-$\alpha$ observations presented in \citetalias{paper1} of this series. We use the results of our stellar wind model to quantify the density, size, and distribution of the planetary exospheres forming as a result of stellar wind sputtering (Section~\ref{sec.sputtering}). {In Section~\ref{sec.observability}, we use a ray tracing technique to predict the observability of the exosphere through transmission spectroscopy in O{\sc i} lines.} Our discussion and conclusions are presented in Section~\ref{sec.conclusions}. | In this paper, we modelled the wind of \hd\ and used it to predict the surface sputter yields and the particle distribution of the {refractory-rich} exosphere for the rocky planets \hd\,b and c. Our stellar wind model is possibly the most well constrained to date after that of the Sun. We used observationally-derived maps of stellar surface magnetic field \citepalias{paper1} for the inner boundary of our 3D wind model. Additionally, the mass-loss rate derived in our wind model ($1.6\e{-14}\msano$) is constrained by Ly-$\alpha$ observations of the stellar astrosphere \citepalias{paper1}. We then used the results of our stellar wind model to quantify wind-induced surface sputtering for the two innermost rocky planets. With that, we were able to estimate the density and structure of the planetary exospheres, on the assumption that both planets do not have strong magnetic fields and have lost both primary (hydrogen-dominated) and secondary (CO$_2$-dominated) atmospheres through escape processes driven by the high-energy stellar flux. Our results can be summarised as follows. The large-scale magnetic field of the planet-hosting star \hd\ can be described as a dipole whose axis is roughly perpendicular to the stellar rotation axis. As a consequence, the stellar wind of \hd\ is highly non-axisymmetric, which implies that planets orbiting in the equatorial plane of the star interact with low and high speed winds in a very short timescale. For example, at every planetary year (3-days orbit), planet `b' interacts with the local stellar wind whose velocities vary from $200$ to $315$~km~s$^{-1}$. A similar level of variation is also seen by planet `c' along its orbit of roughly one week. Because of the close orbital distances, the stellar wind conditions around \hd\,b and c are much harsher than, for example, the solar wind conditions around the Earth, or even around Mercury. If these exoplanets were to have a magnetic field similar to that of the Earth, their magnetospheres would extend out to about 4 planetary radii (only one third of the Earth's magnetospheric size). However, if these planets were to have a magnetic field similar to that of Mercury, their magnetospheres would be crushed into the planets' surface and the stellar wind would directly interact with the planets' crust. In the latter case, due to the close proximity of both planets to the host star, the high {flux of particles} of the incident stellar wind makes the surfaces of the planets to sputter, similarly to what occurs on Mercury. Based on a three-dimensional sputtering model created for Mercury \citep{2015P&SS..115...90P}, we simulated here sputtering processes induced by the stellar wind on the rocky planets \hd\,b and c. Our simulations showed that sputtering processes release refractory elements from the entire dayside surface with velocities sufficiently high to allow for elongated trajectories of the sputtered particles. In particular, we find that oxygen and magnesium are expected to form an extended neutral exosphere with densities larger than 10\,cm$^{-3}$, within several planetary radii. Because of the close proximity of both planets to the host star, a substantial amount of the neutral atoms will quickly be ionised and picked up by the stellar wind. Our simulations suggest the column density of O{\sc i} to be as large as $\sim$10$^{13}$~cm$^{-2}$ close to the day-side of planet `b' and a few times smaller and less extended for planet `c'. The column densities are not symmetric, with enhanced densities ahead of the planets' orbits. This happens due to the angle that the velocity vector of the stellar wind particles makes with the day-side of the planet, when accounting for the orbital motion of the planet through the stellar wind. This enhanced column density ahead of the planet motion could cause an asymmetric transit, with a longer ingress phase than the egress phase, should it be observable. To assess its observability, we used a 3D ray tracing technique to calculate the transit depth in the O{\sc i} $1302.2$~\AA, showing that it is at most $0.042\%$ near line centre, i.e., only a small increase compared to the transit depth in the optical ($0.036\%$). Such a small increase in transit depth in the O{\sc i} $1302.2$~\AA\ line is unlikely to be observable with current UV instrumentation. | 18 | 8 | 1808.00404 |
1808 | 1808.07538_arXiv.txt | We present a formalism for continuum and line emission from random clumpy media together with its application to problems of current interest, including CO spectral lines from ensembles of clouds and radio emission from HII regions, supernovae and star-forming regions. For line emission we find that the effects of clump opacity on observed line ratios can be indistinguishable from variations of intrinsic line strengths, adding to the difficulties in determining abundances from line observations. Our formalism is applicable to arbitrary distributions of cloud properties, provided the cloud volume filling factor is small; numerical simulations show it to hold up to filling factors of \about 10\%. We show that irrespective of the complexity of the cloud ensemble, the radiative effect of clumpiness can be parametrized at each frequency by a single multiplicative correction to the overall optical depth; this multiplier is derived from appropriate averaging over individual cloud properties. Our main finding is that cloud shapes have only a negligible effect on radiation propagation in clumpy media; the results of calculations employing point-like clouds are practically indistinguishable from those for finite-size clouds with arbitrary geometrical shapes. | Radiation propagation in a non-uniform clumpy medium is a common problem in astrophysics. Example continuum applications include IR dust emission from circumnuclear tori in active galactic nuclei \citep[AGN;][]{NENKOVA02, NENKOVA08}, free--free absorption affecting supernova radio light curves and spectra \citep{WEILER04}, radio-millimeter wave thermal emission from single massive stars \citep{IGNACE04} and star-formation induced radio synchrotron emission accompanied by free-free absorption in galaxies \citep{Lacki13}. Spectral line applications include modeling the optical and UV spectra from AGN broad line regions \citep{LAOR06} and interstellar atomic and molecular lines \citep{MHS84, WALL06, WALL07}. A common approach to the analysis of emission from clumpy media, pioneered by \cite{MHS84}, is to assume some geometry for the individual clouds and proceed by averaging over properties along the line of sight (LOS). Some general scaling relations emerged in these works, but remained unexplained. In particular, \cite{IGNACE04} noted that, in their modeling based on spherical clouds, only the distribution in individual cloud optical depths could affect the spectral shape---the cloud radii were irrelevant. In an entirely different approach, \citet{NP84} modeled clumpy media absorption with point-like identical structureless absorbers, characterized by a single property, an optical depth, and no other parameters. Noting that random placement yields a Poisson distribution for the number of absorbers along the LOS, \citeauthor{NP84} derived the effective optical depth of the medium from the mean number of absorbers along the LOS and their common optical depth. \citet{NENKOVA02, NENKOVA08} extended this formalism to the expected emission from a population of such clouds and placed the \citeauthor{NP84} point-like absorbers concept on a more solid footing by showing that the ratio of cloud size to the mean-free-path between clouds is equal to $\phi$, the clouds volume filling factor. Therefore, when $\phi \ll 1$, each cloud appears as a point from its nearest neighbor, thus its geometry can be ignored. Still, the usage of a single optical depth per absorber remained problematic. For example, in the case of a sphere with optical depth $\tau$ along the diameter, the actual optical depth along a LOS can vary from $\tau$, for a LOS through the center, to 0 for a grazing LOS. Here we address this issue, bringing the \citeauthor{NP84} formalism to completion. Starting in \S\ref{se:conttrans} we generalize both clumpy absorption and emission to an arbitrarily complex mixture of clump properties, including variations of these properties along the LOS. The only restrictions are that the medium is random (i.e.\ cloud positions are uncorrelated) and that the propagating radiation does not affect the cloud absorption and emission properties. We investigate via Monte Carlo simulations the range of volume filling factors over which our formalism applies and show that significant departures occur only at relatively large filling factors ($\phi\gtrsim 0.1$; \S\ref{se:simulations}); introduce the concept of a \emph{clumping factor} which modulates the effective opacity of a clumpy medium and depends on the average properties of its clumps (\S\ref{se:K-factor}); and extend the clump formalism to include spectral line absorption (\S\ref{se:specabs}) and emission (\S\ref{se:specemiss}). Section \ref{se:shapes} explores the effects of cloud shapes and shows that they have no significant impact on radiation propagation in a clumpy medium. {In \S\ref{se_examples} we apply our formalism to a couple of current problems involving clumpy emission and absorption by continuum and spectral lines; Appendix \ref{se:ffabs} provides some additional examples}. Section \ref{se:discussion} contains a summary and discussion. | \label{se:discussion} This paper brings to completion the \cite{NP84} approach to radiative transfer in clumpy media, showing that such media can be reliably modeled as collections of structureless clouds (``mega-particles'') characterized by a single property---optical depth. The actual clouds can have a wide range of properties, including different geometrical shapes, opacities, emissivities, spectral shapes, bulk velocities, internal structures and orientations, all of which can vary along the line of sight. With proper averaging, all of these properties can be rigorously encapsulated in an ensemble of identical clouds, and to a good degree of approximation, the geometry of these average clouds is irrelevant.\footnote{\added{It is interesting to note the similarity with approaches taken in the context of radiation propagation in porous media \citep[see][]{Taine08}.}} The simplicity of the formalism presented here has enabled us to readily calculate clumpy emission spectra for a number of current problems, including ensembles of ultra-compact HII regions (\S\ref{se:UCHII}), CO spectral lines (\S\ref{se:line_profiles}) and synchrotron emission accompanied by free-free absorption from supernovae and compact star-bursts in ultra-luminous IR galaxies (Appendix \ref{se:ffabs}). Our results replicate and extend numerous earlier studies. We show that \cite{IGNACE04} correctly identified the unique set of conditions in which clumpy HII regions with a power-law distribution of cloud optical depths can produce $\alpha \simeq 1$ spectra (\S\ref{se:UCHII}). And in the case of spectral line observations we show that it is impossible, even in principle, to distinguish the effect of atomic and molecular abundances on line ratios from the clumping effects of optically thick clouds (\S\ref{se:COlineratio}). The effective optical depth of a clumpy medium (eq.~\ref{eq:contabs}) arises from the result of our formalism for $\<e^{-\tau_\nu}>$, the first moment of the transmission factor distribution. Higher moments can be calculated just as easily---the $m$-th moment, $e^{-m\tau_\nu}$, distribution average is simply $\exp\left[-\N\left(1 - \left< e^{- m\tau_\nu}\right>\right)\right]$, as directly obtained from the derivations in Appendix \ref{se_ap:contabs}. Such moments can yield useful information about the cloud distribution. \citet{TAUBER96} pointed out that a possible route to explore clumpiness is to observe emission lines with very high spectral resolution and signal-to-noise ratios, analyze the fluctuations in brightness temperature present on the line shape, and infer from them the properties of the clumps present in the beam. Assuming identical clumps and employing the \citet{MHS84} model, \citet{TAUBER96} computed the expected fluctuations for a wide range of clump optical depths. Other studies of the fluctuations effect, both earlier \citep{TAUBER91} and later \citep{PIROGOV12}, were restricted to optically thin clumps. Based on our clump formalism it is straightforward to show that the transmission-factor variance obeys \begin{multline} \left\langle\left(e^{-\tau_\nu} - \<e^{-\tau_\nu}>\right)^2\right\rangle = \\ \exp\left[ -\N\left(1 - \left< e^{-2\tau_\nu}\right>\right)\right] - \exp\left[-2\N\left(1 - \left< e^{- \tau_\nu}\right>\right)\right]. \end{multline} This simple expression is completely general and encompasses the results of all previous studies. It enables analysis of spectral variance, as performed by \citet{TAUBER91} and \citet{TAUBER96}, for arbitrary sources without any model restrictions. Such analysis can yield directly \N, the total number of clouds along the LOS. Similar utility exists for higher moments, which can be derived just as easily. While the formalism developed here is quite general, it does rest on some fundamental assumptions. The clump volume filling factor is assumed small enough that departures from Poisson statistics are small. In practice, Monte Carlo simulations (see \S\ref{se:simulations}) show the formalism to give good estimates up to quite large volume filling factors (\about 10\%). When larger filling factors are desired, modeling would have to rely on Monte Carlo simulations. Next, absorption and emission properties of single clouds are assumed unaffected by the radiation generated by the clumpy medium. Relaxing this assumption requires an iterative procedure that starts with initial cloud properties, such as level populations or dust temperature, determined in the absence of cloud emission. In subsequent steps, the clump formalism is used to calculate the expected radiation field, which is then added to the calculation of individual cloud properties and reiterated until convergence. Finally, the formalism assumes random placement in space of individual clouds, such that the presence of another cloud nearby neither increases nor decreases the probability for a cloud at a given position. However, there is evidence suggesting that galactic clouds could be fractal \citep{FALG91, ELM04}, with small high-density clumps embedded within larger lower-density clouds. Nevertheless, there are reasons to expect our formalism to give approximately correct answers even in such a case. Consider a critical size-scale at which the optical depth is approximately unity, such that smaller, denser clumps are optically thick but larger ones are optically thin. Then the fact that the latter are correlated in position with clumps of the critical size has little effect. Smaller, very optically thick cores will be embedded within the already optically thick critical-scale clouds, thus their contributions to the total spectrum will be small. Applying our formalism and considering only the clouds at the critical size and larger should therefore produce reasonably accurate results. We hope to quantify in future work the formalism accuracy when applied to fractal clouds. \appendix | 18 | 8 | 1808.07538 |
1808 | 1808.10487_arXiv.txt | The X-ray source populations within galaxies are typically difficult to identify and classify from X-ray data alone. We are able to break through this barrier by combining deep new Chandra ACIS-I observations with extensive Hubble Space Telescope imaging from the Panchromatic Hubble Andromeda Treasury of the M31 disk. We detect 373 X-ray sources down to 0.35-8.0 keV flux of 10$^{-15}$ erg cm$^{-2}$ s$^{-1}$ over 0.4 square degrees, 170 of which are reported for the first time. We identify optical counterpart candidates for 188 of the 373 sources, after using the HST data to correct the absolute astrometry of our Chandra imaging to 0.1$''$. While 58 of these 188 are associated with point sources potentially in M31, over half (107) of the counterpart candidates are extended background galaxies, 5 are star clusters, 12 are foreground stars, and 6 are supernova remnants. Sources with no clear counterpart candidate are most likely to be undetected background galaxies and low-mass X-ray binaries in M31. The hardest sources in the $1-8$~keV band tend to be matched to background galaxies. The 58 point sources that are not consistent with foreground stars are bright enough that they could be high mass stars in M31; however, all but 8 have optical colors inconsistent with single stars, suggesting that many could be background galaxies or binary counterparts. For point-like counterparts, we examine the star formation history of the surrounding stellar populations to look for a young component that could be associated with a high mass X-ray binary (HMXB). About one third of the point sources are not physically associated with a young population, and are therefore more likely to be background galaxies. For the 40 point-like counterpart candidates associated with young populations, we find that their age distribution has two peaks at 15-20 Myr and 40-50 Myr. If we only consider the 8 counterpart candidates with typical high-mass main sequence optical star colors, their age distribution peaks mimic those of the sample of 40. Finally, we find that intrinsic faintness, and not extinction, is the main limitation for finding further counterpart candidates. | X-ray sources probe the most exotic forms of matter in the universe. Those outside of active galactic nuclei, such as X-ray binaries (XRBs) and supernova remnants (SNRs), can only be detected in nearby galaxies. Chandra and XMM-Newton can resolve hundreds of individual stellar-mass X-ray sources in Local Group galaxies, but outside of the Magellanic Clouds, the identification of counterparts for these stellar mass X-ray sources has been hampered by low spatial resolution X-ray data, difficulty separating background galaxies from stars in optical imaging, and stellar crowding. Over the past decade, our ability to identify high-quality counterpart candidates for X-ray sources outside of the Galaxy and Magellanic Clouds has improved greatly due to the combination of high spatial resolution X-ray imaging with Chandra and resolved stellar photometry with the Hubble Space Telescope (HST). Populations of OB star and background galaxy counterpart candidates have been classified out to distances of 3~Mpc, in particular, M31 \citep[770 kpc;][]{williams2005a, williams2005b,williams2005c,hatzidimitriou2006,williams2014hmxbs}, NGC~300 \citep[2 Mpc;][]{binder2012}, NGC~2403 \citep[3 Mpc;][]{binder2015}, and NGC~404 \citep[3 Mpc;][]{binder2013}. As the nearest massive spiral, M31 has been observed extensively in X-rays. Building on early surveys with the Einstein observatory \citep{vanspeybroeck1979} and ROSAT \citep{supper1997,supper2001}, XMM-Newton has mapped the entire interior of the $D_{25}$ isophotal contour \citep[][hereafter S11]{pietsch2005,stiele2011}, and Chandra has observed the inner disk with the HRC \citep{williams2004}, mapped portions of the disk with ACIS \citep{distefano2004,vulic2016}, and monitored the bulge and nuclear regions for over a decade \citep{kong2002,kaaret2002,garcia2010,li2011}. These surveys have detected dozens of transient X-ray sources, and thousands of persistent sources. Most of the known X-ray sources are unidentified, but many are consistent with emission originating from background active galactic nuclei (AGNs). Others are clearly matched to bright Milky Way foreground stars. However, the most interesting sources are those that may truly be in M31. \citet{pietsch2005} and S11 provided hundreds of source classifications based on variability and hardness ratios, and they identify dozens of SNRs and XRB candidates based on cross-matching with bright stars and star clusters from ground-based imaging and catalogs. With all of this work, only two strong high mass X-ray binary (HMXB) candidates were seen (sources 1579 and 1716 in S11), potentially because of the difficulties of identifying stellar counterparts in the crowded M31 with the spatial resolution available in these data. By comparing the XMM-Newton catalogs and ground-based photometry, \citet{williams2014hmxbs} obtained spectra of dozens of optical counterpart candidates in the M31 field, finding few, if any HMXBs. Most of their spectra showed the counterparts to be background AGN, even though they were targeted to be blue point sources in ground based imaging. While all of this work has significantly advanced our knowledge of M31's X-ray source populations, at this point it remains unclear what fraction of the known X-ray sources actually belong to M31 and which of the sources are background galaxies being viewed through M31. It is also unclear what the nature of most M31 sources is. The very recent HST survey of the northern half of M31, the Panchromatic Hubble Andromeda Treasury \citep[PHAT;][]{dalcanton2012,williams2014}, offers an opportunity to remedy this situation. PHAT is the largest HST mosaic ever assembled, covering a large fraction of the northern M31 disk in 6 HST filters from the near ultraviolet to the near infrared, supplying photometry for over 100 million stars. The high-resolution imaging provides the opportunity to find high-quality counterpart candidates for background galaxies and HMXBs. The resolved photometry allows us to determine the physical characteristics of the stellar populations surrounding the X-ray sources. In addition to allowing us to optically identify background galaxies, we expect the PHAT \citep{dalcanton2012,williams2014} footprint to contain $\gap$20 HMXBs. Measurements of the star formation rate in M31 \citep{williams2003,lewis2015} are $\sim$0.3 M$_{\odot}$ yr$^{-1}$ in the PHAT footprint. The scaling relation between SFR and HMXBs \citep{grimm2003} implies $\sim$20 HMXBs with L$_{\rm X}{>}10^{36}$ erg s$^{-1}$ should be associated with that amount of star formation, and the relation scaling relation in the Magellanic Clouds \citep{antoniou2010} suggests $\sim$100 Be-XRBs, but only a fraction of these ($\sim$20) are expected to have high X-ray luminosities, usually associated with accretion disk systems (i.e., fed by Roche-lobe overflow). Combining the catalogs of \citet{binder2015} with the optical catalogs from the ANGST program \citep{dalcanton2009} suggests a scaling relation between the number of OB stars with M$_V{<}{-}1$ and the number of bright HMXBs (L$_{\rm X}{>}10^{36}$ erg s$^{-1}$). There are $\sim$8$\times$10$^4$ such stars ($m_{f475w}{<}23.55$, $m_{f475w}-m_{f814w}{<}0.5$) in the PHAT footprint, which implies $\sim$30 bright HMXBs. Finally, the summed spectra of sources with L$_{\rm X}{\sim}$(5--10)$\times$10$^{35}$ ergs s$^{-1}$ has a photon index consistent with neutron star HMXBs \citep[see Figure 6 of] []{ShawGreening2009}, consistent with this estimate. So, where are the bright HMXBs in M31? Perhaps they are being missed due to stellar crowding in the M31 disk. To better localize the X-ray sources, we have undertaken a Chandra survey covering much of the PHAT survey area. We designed our survey to provide the largest number of precise positions for the least amount of {\it Chandra} time. Our final observations achieved a 0.35$-$8 keV depth of 3$\times$10$^{-15}$~erg~cm$^{-2}$~s$^{-1}$ (assuming a power-law spectrum with an index of 1.7 and $N_H$=7$\times$10$^{20}$~cm$^{-2}$ (as in, e.g., S11), which corresponds to $\sim$3$\times$10$^{35}$~erg~s$^{-1}$ at the distance of M31 \citep[770 kpc,][]{mcconnachie2005}. This depth should allow us to detect dozens of HMXBs and provide a reliable test of the predicted numbers. Low-mass X-ray binaries (LMXBs) are not as simple to identify, even with {\it HST imaging, as their optical counterparts are too faint to be distinguished. However, based on the stellar mass maps from PHAT \citep{williams2017}, there are $\sim$2$\times$10$^{10}$~M$_{\odot}$ in the region covered by our X-ray data, suggesting a LMXB population with L$_X{>}$3$\times$10$^{35}$ erg s$^{-1}$ of $\sim$100 according to the LMXB X-ray luminosity function (XLF) from \citet{lehmer2014}. By identifying a large fraction of the other sources, we can test this prediction for consistency.} By combining Chandra positions with HST imaging, we simultaneously limit the number of potential counterpart candidates, identify the most likely counterpart to the X-ray source based on HST photometry, and easily resolve many background galaxies, which we expect to dominate the X-ray catalog. For the best X-ray binary candidates, the same data set can be used to constrain the progenitor age, physical characteristics of the secondaries, and provide targets for follow-up optical spectroscopy to measure orbital periods. Ideally, time-resolved spectroscopy of the resulting catalog will ultimately provide clean age and orbital period distributions for a sample of M31 X-ray binaries, which can also be tied to the properties of their local stellar populations. Such a sample will provide quantitative tests for predictions of HMXB production from binary evolution models. In this paper, we present our initial catalog, counterpart candidates, and measure the age distribution of HMXB candidates. In Section 2, we describe the observations of our Chandra survey of the PHAT region, as well as our data reduction technique for measuring the X-ray sources and aligning the Chandra data to PHAT directly. In Section 3, we present our Chandra catalog, cross-matched with the XMM-Newton catalog of S11. We include the most likely optical counterparts in cases where a likely counterpart is present in the HST data. In Section 4, we describe some of the most interesting counterpart candidates, including the best HMXB candidates, and in Section 5 we summarize our work. | We have obtained {\sl Chandra} imaging covering a large fraction of M31 with {\sl Hubble Space Telescope} imaging obtained by the Panchromatic Hubble Andromeda Treasury (PHAT) survey. Combining these data sets, we have produced a catalog of X-ray sources along with their most likely optical counterparts from HST. These optical counterpart candidate identifications allow background galaxies and high-mass X-ray binaries to be separated from other potential hard sources, such as low-mass X-ray binaries. \begin{itemize} \item We find that most counterpart candidates are resolved background galaxies, and that there are over 100 of these, which is consistent with the majority of X-ray sources in the M31 disk field being background contaminants if we assume a similar fraction of the sources with no counterpart candidate are undetected fainter background galaxies. This assumption is consistent both with the expected number of background galaxies estimated from the Chandra Deep Field and with the expected number of LMXBs estimated from the PHAT stellar mass. \item We find about a third of the point source counterpart candidates are not associated with any young stellar populations. \item The number of optical point source candidates (58) is larger than the expected number of bright HMXBs in this region, but it is similar to expectations if about half of the 40 candidates in regions with young stellar populations are indeed HMXBs. \item We find 8 of the point source counterpart candidates have colors typical of single stars, suggesting that many of the point sources in this sample are background galaxies. The number of good HMXB candidates is somewhat below the number expected from the star formation rate and number of OB stars in the region surveyed; however, some of the other point source candidates could be HMXBs with odd colors due to binarity. Further observations will be necessary to determine if M31 actually has as many bright HMXBs as predicted by scaling relations. \item We find that the age distribution of the young populations surrounding the point source counterpart candidates (including the 8 with typical star colors), is peaked at 15-20 Myr and 40-50 Myr in agreement with previous studies in other nearby galaxies \citep{antoniou2010,williams2013,antoniou2016}, but at higher metallicity. \item Based on the extinction results here, dust does not appear to be significantly impeding the searches for optical counterparts. \end{itemize} The production of this catalog is only the beginning of the Chandra-PHAT program. We are currently working to perform and study MCMC fits to the spectral energy distribution from the PHAT photometry of all of the sources within 3-$\sigma$ of X-ray sources using the Bayesian Extinction and Stellar Tool \citep[BEAST][]{gordon2016} similar to those done for the NuSTAR sources in \citet{lazzarini2018}. These fits will likely provide further confirmation of our original classifications presented here, and may result in a few new classifications in cases where there were no obvious candidates in our CMDs or images. Furthermore, the SED fits should allow us to provide physical parameters for the secondaries. Finally, follow-up spectroscopy of our optical point source candidates will help identify more HMXBs, type their secondaries, and measure orbital periods. In turn, we can use all of these measurements in addition to the local star formation histories to place new constraints on X-ray binary formation and evolution models by improving the statistics on their age and mass distributions which the models should reproduce. Support for this work was provided by Chandra Award Number GO5-16085X issued by the Chandra X-ray Observatory Center, which is operated by the Smithsonian Astrophysical Observatory for and on behalf of the National Aeronautics and Space Administration under contract NAS8-03060. M.S.\ acknowledges support by the Deutsche Forschungsgemeinschaft (DFG) through the Heisenberg fellowship SA 2131/3-1 and the Heisenberg professor grant SA 2131/5-1. B.W., M.S. \& P.P. acknowledge partial support for this research through the Chandra Research Visitors Program. | 18 | 8 | 1808.10487 |
1808 | 1808.05151_arXiv.txt | Numerical simulations of pebble dynamics inside gas clumps formed by gravitational instability of protoplanetary discs are presented. We find that dust-mediated Rayleigh-Taylor instabilities transport pebbles inward rapidly via dense metal-rich "fingers". This speeds up sedimentation of small pebbles by up to two orders of magnitude and yet does not impede grain growth because grains of all sizes sediment at the same collective speed as long as Stokes number is less than unity. In simulations with a fixed pebble size, solid planetary cores form if pebble size exceeds a few cm. Pebble growth leads to core formation in some hundreds of years even when pebbles injected into clumps are of mm or smaller sizes. Properties of the gas clump dictate what kind of cores can be made. Low central temperature clumps allow formation of solid cores out of refractory materials, whereas in the highest temperature clumps pebbles of any composition are vaporised and make fuzzy cores only. These results confirm that gravitational instability of protoplanetary discs is a robust mechanism of hatching cores from sub-Earth to Neptune mass, as well as gas giants with massive cores, solid or fuzzy. This mode of planet formation is especially promising for environs too young and distant (such as the ALMA-observed HL Tau disc) or too violent (such as circum-binary planets), to form via the Core Accretion scenarios. | Gas clumps formed by gravitational instability of protoplanetary discs \citep[][]{KratterL16} present a viable environment in which grains can grow, sediment to the centre and form a solid core there \citep{Kuiper51,Kuiper51b,McCreaWilliams65,CameronEtal82,Boss98}. The central temperatures of these clumps vary, depending on their mass and evolutionary state, from $\sim 100$~K to $\sim 2000$~K \citep{Bodenheimer74}, in principle allowing grains of various compositions to reach the clump centre. \cite{Kuiper51b} believed that planets form on fixed orbits. We now know that massive self-gravitating gas discs hatch clumps by disc fragmentation at separations $\sim 100$~AU \citep{Gammie01,Rafikov05} but the clumps may migrate closer to the star in a matter of a few thousand years \citep{MayerEtal04,VB06,VB10,MachidaEtal10,BaruteauEtal11,MichaelEtal11,MachidaEtal11}. Those clumps that manage to contract and collapse into second cores \citep[also called post-collapse gas giants, or "hot start models" in different contexts][]{Larson69,BurrowsEtal00}, and survive the migration phase, become gas giant planets. Clumps that contract too slowly are tidally disrupted \citep{BoleyEtal10,Nayakshin10c}. If a solid core is synthesized inside the clump by the time it is disrupted, the core is released back into the world. However, the total condensible mass of metals inside a gas clump of mass $M$ is only $\sim 3 \mearth (M/1 M_J) (Z/0.01)$, where $Z$ is clump metallicity. Grains also need to be as large as 1 cm in radius for an efficient grain sedimentation into the core, and this may not occur in time before the clump is disrupted. Finally, the internal regions of the clump may be too hot for grains to exist. Previous {\em isolated} clump studies found that these challenges are not easily overcome \citep{HelledEtal08,HS08,ForganRice13b}. However, the dust content of gas clumps may be far greater due to accretion of $\sim 1$ mm or larger grains from the disc \citep{HN18} via a process known as pebble accretion \citep{OrmelKlahr10,JohansenLacerda10,LambrechtsJ12,LambrechtsEtal14}. Furthermore, spiral arms and gas clumps may be enriched with solids already at birth \citep{RiceEtal04,BoleyDurisen10,BoleyEtal11a,GibbonsEtal12,GibbonsEtal14}. \cite{NayakshinFletcher15,Nayakshin16a} included the process of pebble accretion in their population synthesis, with results showing some promise in terms of core masses, compositions, orbital separations and host star metallicity correlations \citep[for a broad comparison of the model results with observations, see][]{Nayakshin_Review}. However, the processes of grain sedimentation and core formation inside gas clumps were studied previously either analytically or via 1D spherically symmetric numerical codes only. In this paper we present first 3D numerical simulations of coupled gas and dust dynamics inside the gas clumps. We are in particular interested in the fate of the additional grains accreted by the clump from the parent disc because these grains may outnumber by total mass those native to the clump. To achieve higher numerical resolution, isolated gas clumps are studied here but the initial conditions are tailored to mimic clumps in their protoplanetary disc birth environment. We start with simulations in which grain size is fixed, the initial conditions are spherically symmetric, but later relax these assumptions. Table 1 (see \S \ref{sec:ic}) gives a summary of simulations presented here and main results learned from these. Animations of two simulations, Sp1Z1a01F and DarkCollapse, are available via online supplementary material, and at these links, respectively: \noindent {\scriptsize \verb!https://www.dropbox.com/s/7e56pxlnhtkrqk6/SpZ1a01N8e5.mp4?dl=0!} {\scriptsize \verb!https://www.dropbox.com/s/wvnailpusfg8x7t/DarkCollapse.mp4?dl=0!} | Simulations presented here show that clumps that accrete pebbles from their parent discs can make high-Z cores more rapidly than assumed based on earlier closed-box 1D models of the clumps \citep[e.g.,][]{HS08,BoleyEtal10,Nayakshin10c,ForganRice13b}. In general, pebbles loaded onto the clump sediment through the outer envelope rapidly in the test particle regime, and then stall in higher density regions. The dust Rayleigh-Taylor instability then develops, transporting them in in a matter of tens to hundreds of years. Small grains grow and large grains fragment in the metal enriched central part of the clump, also on time scales of tens to hundreds of years. A few cm or larger sized pebbles then get locked into solid cores. The outcome of these processes depends on pebble composition, how hot the centre of the clump is, and how long it has to live before its tidal disruption. Although these external factors are not modeled in this paper, previous 1D models of planet formation by gravitational instability that {included both dust physics and pebble accretion} \citep[e.g.,][]{NayakshinFletcher15} compare with many observational facts favorably \citep[for details see \S 9 in][]{Nayakshin_Review}. 3D simulations presented here however show that cores can be made even more rapidly due to dust-RT instability, and that even small pebbles loaded into the clump tend to concentrate into the clump central regions rather than be spread around the clump uniformly \citep[as was assumed in][]{Nayakshin15a}. As a result, we found that gravitational instability clumps may form gas giants with fuzzy cores if the central regions of the clump are hotter than $\sim 1500$~K. This may be relevant to the recent {\em Juno} satellite Jupiter's gravity measurements that indicate that its core may be fuzzy rather than solid \citep{WahlEtal17}. Formation of cores inside the gaseous clumps formed by gravitational instability is a promising and probably required mechanism to explain planets found in circumstances unfavorable to formation by Core Accretion \citep{PollackEtal96}. For example, the suspected $\sim$ Neptune mass planets in the $\sim 1$~Myr-old young disc of HL Tau \citep{BroganEtal15,DipierroEtal15} should have formed after $\sim 10^8$ years in the classical Core Accretion scenario \citep[e.g.,][]{KB15}. These low mass planets could not form by the pure gas disc fragmentation as such objects are at least $\sim 1 \mj$ in mass \citep{BoleyEtal10}. Another promising application of the theory is close circum-binary planets, where binary kicks lead to very violent planetesimal-splitting collisions \citep{LinesEtal14}. In the context of gravitational instability, these planets could have formed inside the protective envelope of the self-gravitating gas clump, initially at large separation from the binary centre. When such a clump migrates close enough to be disrupted, its ready-made-core or planet could be safely deposited onto a much smaller orbit. Our simulations however do not include radiative cooling of the clumps and feedback from growing massive cores \citep{Nayakshin16a}. These effects may dictate the resulting metallicity correlations of objects made by gravitational instability, from planetary debris and sub-Neptune planets \citep{FletcherNayakshin16a} to massive planets and brown dwarfs \citep{NayakshinFletcher15}. 3D simulations including these effects shall be reported elsewhere. Finally, note that the instability presented here is probably related to the "drafting instability" recently found by \cite{LambrechtsEtal16} in the context of vertical grain settling in protoplanetary discs. The authors also predicted existence of this instability in the envelopes of growing gas giant planets in the context of the Core Accretion model for planet formation. Our results therefore echo their funding for planets formed via gravitational disc instability. | 18 | 8 | 1808.05151 |
1808 | 1808.09978_arXiv.txt | We study the properties of 30 spectroscopically-identified pairs of galaxies observed during the peak epoch of star formation in the universe. These systems are drawn from the MOSFIRE Deep Evolution Field (MOSDEF) Survey at $1.4 \leq z \leq 3.8$, and are interpreted as early-stage galaxy mergers. Galaxy pairs in our sample are identified as two objects whose spectra were collected on the same Keck/MOSFIRE spectroscopic slit. Accordingly, all pairs in the sample have projected separations $R_{{\rm proj}}\leq 60$~kpc. The velocity separation for pairs was required to be $\Delta v \leq 500 \mbox{ km s}^{-1}$, which is a standard threshold for defining interacting galaxy pairs at low redshift. Stellar mass ratios in our sample range from 1.1 to 550, with 12 ratios closer than or equal to 3:1, the common definition of a ``major merger." Studies of merging pairs in the local universe indicate an enhancement in star-formation activity and deficit in gas-phase oxygen abundance relative to isolated galaxies of the same mass. We compare the MOSDEF pairs sample to a control sample of isolated galaxies at the same redshift, finding no measurable SFR enhancement or metallicity deficit at fixed stellar mass for the pairs sample. The lack of significant difference between the average properties of pairs and control samples appears in contrast to results from low-redshift studies, although the small sample size and lower signal-to-noise of the high-redshift data limit definitive conclusions on redshift evolution. These results are consistent with some theoretical works suggesting a reduced differential effect of pre-coalescence mergers on galaxy properties at high redshift -- specifically that pre-coalescence mergers do not drive strong starbursts. | \label{sec:intro} Galaxies grow in mass through a combination of mergers with other galaxies and smooth accretion of baryons and dark matter. Predicting the frequency of both major (i.e., with roughly equal masses ) and minor (i.e., with significantly unequal masses) mergers as a function of galaxy mass and redshift is therefore an important component of hierarchical models of structure formation \citep[e.g.,][]{hopkins2010}. At the same time, obtaining empirical constraints on such merger rates as a function of galaxy mass and redshift represents a key goal for observations of galaxy evolution \citep[e.g.,][]{lotz2011}. In addition to quantifying merger rates, both models and observations aim to describe the impact of galaxy interactions on the properties of merging and coalesced galaxies. Simulations of star-forming galaxy mergers predict a characteristic progression of the star-formation rate (SFR) throughout the merger event. Relative to the time prior to the merger, the SFRs of the merging galaxies are elevated during their extended gravitational interaction, and ultimately peak when the galaxies coalesce \citep[e.g.,][]{mihos1996,hopkins2008,cox2008}. The degree of enhancement in SFR is predicted to depend on galaxy mass ratio. For example, \citet{cox2008} has demonstrated that mergers with mass ratio smaller than 3:1 lead to much stronger bursts of star formation than mergers with larger mass ratios. Additional factors affect the strength of the merger-induced starburst, such as the orientation of the orbits of merging galaxies, as well as their structural properties and gas fractions. In the local universe, the most luminous systems, (i.e., ultraluminous infrared galaxies; ULIRGs), appear to be dominated by advanced-stage major mergers during or just after coalescence \citep[e.g.][]{sanders1996,tacconi2002}. Pre-coalescence stages of merging at $z\sim 0$ have been traced by galaxy pairs. The Sloan Digital Sky Survey (SDSS) has yielded a statistical sample of such pairs \citep[e.g.,][]{ellison2008, scudder2012, scudder2015,patton2011,patton2013}, identified as galaxies separated by both a small projected radius (with upper limits on $R_{{\rm proj}}$ ranging from 30 to 80~kpc) and small radial velocity difference (with upper limits on $\Delta v$ ranging from $200$ to $500\mbox{ km s}^{-1}$). Members of these galaxy pairs are characterized by both enhanced SFRs \citep[e.g., $\sim 60$\% out to 30~kpc;][]{scudder2012} and depressed gas-phase oxygen abundances \citep[e.g., $\sim 0.02$~dex;][]{scudder2012} relative to a control sample of isolated galaxies matched in stellar mass. Such differences are consistent with theoretical models of galaxy mergers in which an increase in SFR accompanies the inflow of gas into the central regions of the merging galaxies, which also tends to dilute the metal content of the interstellar medium \citep[ISM;][]{hopkins2008,bustamante2018}. A similar enhancement in SFR in galaxy pairs has been detected out to $z\sim 1$ \citep{lin2007,wong2011}. \begin{figure*}[t!] \includegraphics[width=\textwidth]{f1.pdf} \caption{Gallery of two-dimensional spectra of galaxy pairs at $1.37 \leq z \leq 1.70$. Two panels are shown for each galaxy pair, which zoom in on the rest-frame wavelength ranges centered on [OIII]$\lambda 5007$ and H$\alpha$ for the primary target, and correspond to $\Delta z=\pm 0.01$. At these redshifts [OIII]$\lambda 5007$ falls in the observed $J$ band, while H$\alpha$ falls in the observed $H$ band. For each pair, emission from the primary target galaxy is circled in red, while that from the serendipitous companion is circled in blue. 3D-HST v4.1 catalog numbers \citep{skelton2014} are given to the left of the [OIII]$\lambda 5007$ panel, with red and blue color-coding corresponding to the circles. Spectral cut-outs are scaled to the same vertical size on the page for display purposes, resulting in a variable angular scale in that dimension. Accordingly, we provide a white vertical scale bar in each panel indicating the extent of 30 proper kpc, which corresponds to 3\secpoint65 at the median redshift of the sample.} \centering \label{fig:z12d} \end{figure*} \begin{figure*}[t!] \includegraphics[width=0.95\textwidth]{f2.pdf} \caption{Gallery of two-dimensional spectra of galaxy pairs at $1.90 \leq z \leq 2.70$. Two panels are shown for each galaxy pair, which zoom in on the rest-frame wavelength ranges centered on [OIII]$\lambda 5007$ and H$\alpha$ for the primary target, and correspond to $\Delta z=\pm 0.01$. At these redshifts [OIII]$\lambda 5007$ falls in the observed $H$ band, while H$\alpha$ falls in the observed $K$ band. All symbols, labels, and scale bars as in Figure~\ref{fig:z12d}. The $z=2.3$ pair GOODSN-24825/25017 lacks coverage of H$\alpha$ because it was observed as a filler target on a MOSDEF mask targeting $1.37\leq z \leq 1.70$ galaxies and therefore not observed in the $K$ band. } \centering \label{fig:z22d} \end{figure*} \begin{figure*}[t!] \includegraphics[width=\textwidth]{f3.pdf} \caption{Gallery of two-dimensional spectra of galaxy pairs at $2.95 \leq z \leq 3.80$. One panel is shown for each galaxy pair, which zooms in on the rest-frame wavelength ranges centered on [OIII]$\lambda 5007$ for the primary target, and corresponds to $\Delta z=\pm 0.01$. At these redshifts [OIII]$\lambda 5007$ falls in the observed $K$ band. All symbols, labels, and scale bars as in Figure~\ref{fig:z12d}. } \centering \label{fig:z32d} \end{figure*} Galaxy mergers have now been identified out to $z\sim 6$ \citep{ventou2017}. In the early universe (i.e., at $z>1$), primarily one of two techniques is employed to flag merging systems. First, it is common to use morphological signatures to identify ongoing or recently completed merger events. Galaxies have been visually classified as mergers on the basis of morphological features such as tidal tails and bridges, and double nuclei \citep{lofthouse2017}, and also identified as interacting based on non-parametric morphological statistics such as the Gini and $M_{20}$ coefficients \citep{lotz2004}, or the concentration (C), asymmetry (A), and clumpiness (S) statistics \citep{conselice2014}. The second technique for flagging mergers is through galaxy pairs. Many studies aiming to quantify the merger fraction and rate at $z>1$ have been based on photometric pairs, which consist of galaxies within a small projected radius and small difference in photometric redshift \citep[e.g.,][]{man2012,man2016,mantha2018,williams2011}. In some cases \citep{bluck2009}, the photometric redshift for only one of the galaxies is known. The possibility of contamination by chance projections must therefore be accounted for, especially when the redshift of a potential companion galaxy is unknown. Recently, merging pairs at $z>1$ have also been identified spectroscopically, based on rest-frame ultraviolet spectra \citep{tasca2014,ventou2017}. However, the sensitivity of rest-frame UV features to large-scale galaxy outflows \citep{pettini2001,shapley2003,steidel2010} limits the accuracy with which such galaxy systemic redshifts, and therefore merger dynamics, can be measured. To date, most studies of merging pairs at $z>1$ have focused on global statistics such as the merger fraction and rate, as opposed to systematic studies of the impact of close interactions on the properties of merging galaxies. In this work, we focus on the latter, based on a sample of galaxy mergers identified at $1.4 \leq z \leq 3.8$ within the MOSFIRE Deep Evolution Field (MOSDEF) survey \citep{kriek2015}. The extensive rest-optical spectroscopic coverage of the MOSDEF survey enables us to assemble a clean sample of galaxy pairs that are not only close on the sky but also in redshift space. With spectroscopic pairs, there is little possibility of contamination by chance projections of completely unassociated galaxies. Such chance projections can arise when pairs are identified on the basis of photometric redshifts, given their associated uncertainties at high redshift.\footnote{We note that proximity in redshift space does not guarantee merging, as galaxy pairs offset by tens of proper kpc in $R_{proj}$ and with line-of-sight velocity separations of up a few hundreds of $\mbox{ km s}^{-1}$ may not be bound and destined to merge \citep{moreno2013}. However, with spectroscopic redshift measurements, we can at least apply the same proximity criteria in velocity space that is used for studies of local galaxy pairs.} Furthermore, we have estimated key galaxy properties such as SFR, stellar mass ($M_*$), and gas-phase oxygen abundance for both merging and isolated systems, and can therefore study for the first time the effect of interactions on star-formation activity and chemical enrichment in distant star-forming galaxies. In Section~\ref{sec:mosdef}, we present the details of the MOSDEF survey and the galaxy properties analyzed in this work. Section~\ref{sec:pairs} discusses the selection and properties of spectroscopically-determined merging pairs in MOSDEF, while Section~\ref{sec:control} describes the selection of our control sample of isolated galaxies used for systematic comparison with mergers. In Section~\ref{sec:properties}, we investigate the effect of mergers on star formation and metal enrichment through analysis of the SFR-$M_*$ main sequence and stellar mass-metallicity relation (MZR) for both merging pairs and isolated control galaxies. We present a discussion of our results and describe future work in Section \ref{sec:discussion}. Throughout this paper, we adopt cosmological parameters of $H_0=70 \mbox{ km s}^{-1} \mbox{ Mpc}^{-1}$, $\Omega_M = 0.30$, and $\Omega_{\Lambda}=0.7$. | \label{sec:discussion} We have demonstrated that merging galaxy pairs at $1.4 \leq z \leq 2.6$ are not characterized by elevated SFRs or significantly diluted metallicities relative to isolated systems of the same stellar mass. Although we will require a larger pairs sample to place these results on a more secure statistical footing, it is worthwhile to consider the implications of the suggested trends, which run contrary to the patterns uncovered among interacting galaxies at lower redshift. As we discuss below, there is a theoretical basis for the apparent evolution in the differential star-forming properties of merging systems. Here we compare our results with other differential studies of merging pairs at $z>1$, consider the predictions from simulations, and discuss evidence for AGN activity in our merging systems. We conclude by listing some promising future directions for the study of interacting galaxies at high redshift. \subsection{Comparisons with Other Observational Work}\label{sec:discussion-comparison} Our study represents the first controlled differential comparison of SFRs in interacting and isolated galaxies at $z>1$. However, \citet{divoy2014} previously considered the question of the relationship between metallicity and small-scale environment. These authors analyzed a sample of 49 star-forming galaxies at $z\sim 0.9-1.8$ with VLT/SINFONI integral field unit emission-line maps and associated $N2$ metallicity measurements. In this sample, 12 systems are identified as ``interacting" based on the presence of a companion within $\Delta v=500 \mbox{ km s}^{-1}$ and $R_{{\rm proj}}=30 h^{-1}$~kpc \citep{lopezsanjuan2013}. \citet{divoy2014} find a measurable depression in median metallicity of at least 0.13~dex for their ``interacting" sample relative to a control sample of 37 isolated galaxies at similar median stellar mass. This depression is significantly larger than what is observed in the local universe \citep[$\leq 0.05$~dex;][]{ellison2008,scudder2012}, and in contrast to the lack of significant offset in metallicity that we find for our pairs relative to the control sample. In order to obtain a more robust result at high redshift, we require a significantly larger sample of galaxy pairs with metallicity measurements, and deeper spectroscopy sufficient to detect H$\beta$ and [NII]$\lambda 6584$ for a larger fraction of the sample in order to reduce the number of metallicity upper limits. \subsection{Expectations From Simulations}\label{sec:discussion-simulations} The lack of enhancement in SFR observed for the MOSDEF pairs sample is consistent with recent predictions from numerical simulations of galaxy formation. \citet{fensch2017} run a suite of pc-scale galaxy merger simulations, representing low-redshift galaxies with gas fractions of 10\%, while $z=2$ galaxies are simulated with gas fractions of 60\%. For the same orbital parameters, it is found that the gas-poor merger simulations approximating local galaxies feature boosts in the SFR of the merging galaxies of an order of magnitude or more over an extended period of hundreds of Myr prior to coalescence, while the gas-rich simulations approximating $z=2$ show only mild increases in SFR, and only at coalescence. Based on a suite of binary galaxy merger simulations with identical orbital parameters but nine different initial gas fractions ranging from $M_{gas}/M_* = 0.04$ to $1.78$, \citet{scudder2015} similarly find that the enhancement in SFR during the merger is anti-correlated with initial gas fraction. \citet{fensch2017} attribute the difference in the evolution of the SFR during the merger to differences in the increase of both central gas inflows and compressive turbulence in the ISM during mergers at $z=2$ and $z=0$. Both gas-rich ($z=2$) and gas-poor ($z=0$) merger simulations reach similar peak central gas inflow rates fueling star formation, but, given that the gas-rich simulations start off with pre-merger baseline gas inflow rates that are an order of magnitude higher than those in the gas-poor simulations, the enhancement in gas inflow and corresponding SFR is significantly weaker. \citet{fensch2017} then attribute the lack of increase in ISM turbulence during $z=2$ mergers to both an ISM velocity dispersion that is higher pre-merger and harder to additionally stir-up \citep{wisnioski2015}, and a clumpy ISM architecture with an associated tidal field that also suppresses an increase in turbulence \citep{genzel2008}. These simulations do not incorporate cosmological accretion or the more compact nature of $z=2$ galaxies relative to those at $z=0$ of the same mass, though the authors argue that differences in gas fraction and ISM structure are the dominant effects underlying the observed difference in SFR evolution for high-redshift mergers. Analysis of the lower-resolution Horizon-AGN cosmological hydrodynamical simulation by \citet{martin2017} also predicts that the enhancement in SFR at fixed mass due to mergers is most pronounced at low redshift, and undetectable at $z>1.5$, when the ambient pre-merger level of star formation is an order of magnitude higher than in the local universe. There is also recent theoretical work focusing on the evolution of ISM metallicity during galaxy mergers. Based on a sample of 70 gas-rich mergers traced at $z<1.5$ in the Auriga cosmological simulation, \citet{bustamante2018} find that a period of metallicity dilution typically occurs during merger events, reaching a magnitude of $\Delta Z=-0.1$~dex for major mergers at projected separations $<10$~kpc, and at least a few hundreths of a dex depression for both major and minor mergers at separations of $<30$~kpc. These results are roughly consistent with observations of $z\sim 0$ merging pairs by \citet[e.g.,][]{scudder2012}. However, it is also worth noting that the pre-coalescence depressions in metallicity recovered in these simulations are consistent with those inferred from the fundamental metallicity relation \citep[FMR;][]{mannucci2010,ellison2008fmr}. Specifically, given the relation among $M_*$, SFR, and gas-phase oxygen abundance in the FMR, the metallicity during the mergers is consistent with predictions from the FMR, given the evolution in SFR and $M_*$ of the merging galaxies. Accordingly, given that our merging pairs show no offset from the control sample in the SFR vs. $M_*$ relation, and given that these pairs represent a pre-coalescence phase, the expectation from the simulation results of \citet{bustamante2018} is that pairs in MOSDEF should simply follow the same relationship among $M_*$, SFR, and metallicity that is inferred at $z\sim 2$ for the MOSDEF sample as a whole \citep{sanders2018}. Furthermore, given the size of the error bars on our median stacked metallicities, we are not sensitive to differences of $\leq 0.05$~dex as observed by \citet{ellison2008} and \citet{scudder2012}, and predicted for all but the most extreme merger events in \citet{bustamante2018}. Finally, we note that \citet{torrey2017} have also found evidence for metallicity dilution during mergers in the IllustrisTNG cosmological simulation, but the detailed example analyzed in that work consisted of a merger at $z\sim 0$. It is not clear how the analogous results would differ at $z\sim 2$. High-resolution ``$z=2$" simulations like those of \citet{fensch2017} are required to address the question of metallicity dilution in high-redshift mergers. For robust conclusions, such simulations must also include cosmological gas accretion and track metal enrichment as well as star formation. \subsection{The AGN Fraction in MOSDEF Pairs}\label{sec:discussion-agn} Our samples of galaxies in pairs are indistinguishable from their isolated counterparts in terms of star formation and metallicity at fixed stellar mass. The extensive multi-wavelength data available for MOSDEF targets also enables AGN classifications on the basis of X-ray luminosity, infrared colors, and/or rest-optical emission line ratios \citep{coil2015,azadi2017}. Given the connection between galaxy mergers and black hole fueling inferred from numerical simulations \citep[e.g.,][]{hopkins2006,dimatteo2005}, much observational work has been devoted to exploring the link between galaxy interactions and AGN activity, with mixed results. For example, \citet{ellison2011} demonstrate an enhancement in AGN fraction among galaxy close pairs at $z\sim 0$, with up to a factor of 2.5 enhancement in AGN fraction for pairs with projected separations less than 40~kpc. Along the same lines, \citet{goulding2018} use machine-learning classifications of merging systems at $z\sim 0-0.9$, finding an AGN fraction (as estimated by WISE infrared colors) elevated by a factor of $\sim 2-7$ among interacting relative to non-interacting systems. Many other investigations have focused instead on the merger fraction among AGNs (as opposed to the AGN fraction among mergers), finding that AGNs do not show a significantly enhanced fraction of mergers relative to their inactive counterparts \citep[e.g.,][]{kocevski2012,gabor2009,hewlett2017}. Exploring the merger fraction among AGNs and non-AGNs provides constraints on how important interactions are for driving AGN activity relative to other more secular processes \citep{bournaud2011}. Given the design of our experiment and the incompleteness of our pairs sample (see Section~\ref{sec:control}), we are not in the position to quantify the differential merger fraction among AGNs and non-AGNs. However, we can obtain complete estimates of the AGN fraction among merging and isolated galaxies, which indicates how likely it is for mergers to trigger AGN activity. Based on the X-ray, infrared, and rest-optical emission-line AGN classifications for MOSDEF galaxies, we find that in our full pairs sample the fraction of galaxies satisfying at least one of these AGN criteria is 10 out of 59 ($16.9\% \pm 5.3$\%, with the error based on simple Poisson statistics). Focusing on the $1.4 \leq z \leq 2.6$ subsample for which we compared SFRs and metallicities with their isolated counterparts, we find 9 out of 47 ($19.1\% \pm 6.3$\%) are classified as AGNs. We also note that 7 out of 36 ($19.4\% \pm 7.3$\%) of these systems with $R_{{\rm proj}}\leq 30$~kpc -- i.e., the same fraction -- satisfy at least one of the AGN criteria. The corresponding AGN fractions for the full and $1.4 \leq z \leq 2.6$ control subsamples are 42 out of 372 ($11.3\% \pm1.7$\%) and 32 out of 242 ($13.2\%\pm2.3$\%). Accordingly, we find a higher AGN fraction among pairs relative to controls, with a factor of at least $\sim 1.5$ enhancement in the pair AGN fraction at $1.4 \leq z \leq 2.6$. However, given the small sample size of the pairs, this difference is not highly significant. Increasing the sample size of pairs and control galaxies by an order of magnitude will enable robust statistics on the differential AGN fraction between pairs and control galaxies. \subsection{Future Work}\label{sec:discussion-future} Our analysis comprises an early step in characterizing the properties of merging galaxies at high redshift. With only 30 pairs in total, 24 of which we analyze in detail along with a carefully-defined control sample of isolated galaxies, the statistical power of our sample is limited. For example, while \citet{ellison2008} and \citet{scudder2012} divided their samples of almost 2000 merging pairs into multiple bins of projected separation and mass ratio to explore how the physical properties of merging galaxies depend on each of these merger characteristics, our current small sample size precludes such division. In addition to obtaining a significantly larger sample of merging pairs, we must also obtain significantly deeper rest-optical emission-line spectra, so that the fraction of limits in SFR(H$\alpha$) and $12+\log(\mbox{O/H})$ is reduced to a negligible contribution across a sufficiently wide mass range, and we can therefore analyze the distributions of individual merging and isolated galaxies in the spaces of SFR(H$\alpha$) vs. $M_*$ and $12+\log(\mbox{O/H})$ vs. $M_*$. We also need to apply an effective technique for identifying later-stage mergers closer to coalescence, when the effects of SFR enhancement and metallicity dilution are predicted to be strongest \citep{bustamante2018}. At low redshift, non-parametric morphological measures are commonly used to identify such mergers \citep{lotz2004}, but \citet{abruzzo2018} demonstrate with mock observations of cosmological simulations that these same merger morphological classifications at $z\sim 2$ are significantly contaminated by non-merging galaxies and should not be applied to {\it HST} images of high-redshift galaxies. Another route to studying later-stage mergers would consist of building up a significantly larger (i.e., two orders of magnitude) sample of spectroscopically-confirmed pairs with $R_{{\rm proj}}<10$~kpc pairs and $\Delta v \leq 500 \mbox{ km s}^{-1}$. With larger galaxy pairs samples accompanied by complete sets of emission-line detections and a larger range of merger stages, we will truly be able to test models of galaxy mergers in the early universe. \begin{deluxetable*}{lrccrccrrr} \tablecolumns{8} \tablewidth{7in} \tablecaption{Observed Properties of MOSDEF Pairs\label{tab:basic}} \tablehead{ \colhead{Field} & \colhead{ID\tablenotemark{a}} & \colhead{$z_{\rm{MOSFIRE}}$\tablenotemark{b}} & \colhead{$m_{\rm{F160W}}$\tablenotemark{c}} & \colhead{ID\tablenotemark{a}}& \colhead{$z_{\rm{MOSFIRE}}$\tablenotemark{b}} & \colhead{$m_{\rm{F160W}}$\tablenotemark{c}} &\colhead{ $R_{\rm{proj}}$\tablenotemark{d}} & \colhead{$\Delta v$\tablenotemark{e}} & \colhead{$M_{*{\rm Larger}}/M_{*{\rm Smaller}}$\tablenotemark{f}}\\ \colhead{} & \colhead{(Primary)} & \colhead{} & \colhead{} & \colhead{(Serendip)} & \colhead{} & \colhead{} &\colhead{(kpc)} & \colhead{($\mbox{km s}^{-1}$)} & \\ } \startdata \\ \multicolumn{9}{c}{$1.37\leq z \leq 1.70$} \\ \hline COSMOS & 8381 & 1.4049 & 21.21 & 8490 & 1.4052 & 23.55 & 38.4 & 37 & 7.6 \\ COSMOS & 19325 & 1.6013 & 23.86 & 19292 & 1.6007 & 24.16 & 11.1 & 74 & 1.9 \\ GOODSN & 10044 & 1.6006 & 23.14 & 10041 & 1.6039 & 25.33 & 53.1 & 380 & 11.0 \\ AEGIS & 35056 & 1.6534 & 23.81 & 35075 & 1.6530 & 24.27 & 5.7 & 38 & 1.1 \\ AEGIS & 39567 & 1.5805 & 21.95 & 39897 & 1.5806 & 23.89 & 10.7 & 17 & 7.8 \\ AEGIS & 16339 & 1.5291 & 21.13 & 16026 & 1.5320 & 21.69 & 4.8 & 338 & 1.8 \\ AEGIS & 3237 & 1.6667 & 22.38 & 3478 & 1.6657 & 22.66 & 6.8 & 113 & 8.1 \\ \hline\\ \multicolumn{9}{c}{$1.90\leq z \leq 2.61$} \\ \hline COSMOS & 1740 & 2.2999 & 23.89 & 1795 & 2.2985 & 26.04 & 39.4 & 121 & 4.4 \\ COSMOS & 10719 & 2.2465 & 23.61 & 10766 & 2.2505 & 25.07 & 9.9 & 368 & 6.8 \\ COSMOS & 471 & 2.0731 & 23.52 & 532 & 2.0678 & 24.55 & 52.5 & 521 & 3.8 \\ COSMOS & 21492 & 2.4721 & 20.99 & 21598 & 2.4786 & 24.06 & 20.5 & 561 & 7.8 \\ COSMOS & 14849 & 1.9265 & 22.08 & 14968 & 1.9233 & 21.99 & 11.8 & 330 & 3.5 \\ COSMOS & 25229 & 2.1813 & 22.79 & $-$9999\tablenotemark{g} & 2.1826 & \nodata\tablenotemark{g} & 12.8 & 130 & \nodata \\ COSMOS & 7433 & 2.1662 & 24.03 & 7417 & 2.1661 & 24.70 & 8.0 & 8 & 2.0 \\ GOODSN & 17748 & 2.2325 & 24.20 & 17714 & 2.2349 & 23.67 & 10.2 & 221 & 1.5 \\ GOODSN & 24825 & 2.3347 & 23.63 & 25017 & 2.3359 & 23.47 & 29.7 & 110 & 1.3 \\ GOODSN & 23344\tablenotemark{h} & 2.4839 & 23.25 & 23418 & 2.4889 & 26.39 & 45.1 & 433 & 550.0 \\ GOODSN & 23344\tablenotemark{h} & 2.4839 & 23.25 & 23339 & 2.4828 & 23.53 & 12.5 & 92 & 7.6 \\ GOODSN & 23869 & 2.2438 & 23.15 & 24074 & 2.2433 & 24.37 & 50.2 & 46 & 5.0 \\ GOODSN & 12302 & 2.2756 & 21.97 & 12172 & 2.2754 & 25.06 & 10.5 & 16 & 17.0 \\ AEGIS & 18543 & 2.1387 & 23.15 & 18454 & 2.1384 & 24.23 & 12.4 & 33 & 2.0 \\ AEGIS & 31108 & 2.3547 & 22.68 & 31317 & 2.3570 & 23.30 & 24.3 & 205 & 5.4 \\ AEGIS & 36050 & 2.5285 & 23.31 & 36180 & 2.5325 & 23.07 & 17.7 & 339 & 3.9 \\ AEGIS & 29114 & 2.3519 & 24.46 & 29045 & 2.3500 & 24.94 & 8.6 & 175 & 2.3 \\ \hline\\ \multicolumn{9}{c}{$2.95\leq z \leq 3.80$} \\ \hline COSMOS & 23183 & 3.1211 & 24.45 & 23192 & 3.1227 & 24.87 & 9.2 & 123 & 1.6 \\ COSMOS & 24579 & 3.2521 & 25.04 & 24596 & 3.2566 & 25.84 & 11.6 & 315 & 4.4 \\ COSMOS & 2360 & 3.0338 & 24.34 & 2344 & 3.0344 & 25.08 & 16.5 & 39 & 2.2 \\ GOODSN & 28202 & 3.2325 & 23.87 & 28209 & 3.2336 & 23.78 & 5.4 & 76 & 1.1 \\ GOODSN & 15694 & 3.3233 & 24.71 & 15566 & 3.3203 & 26.68 & 13.1 & 207 & 9.5 \\ AEGIS & 30847 & 3.4349 & 23.51 & 30691 & 3.4431 & 24.92 & 16.9 & 553 & 3.1 \enddata \tablenotetext{a}{Galaxy ID in the 3D-HST v4.1 catalogs \citep{momcheva2016}.} \tablenotetext{b}{Emission-line redshift measured from MOSFIRE spectroscopy as part of the MOSDEF survey.} \tablenotetext{c}{{\it HST}/WFC3 F160W magnitude on the AB system.} \tablenotetext{d}{Projected separation in kpc, based on the angular separation between the pair galaxies and the redshift of the primary galaxy.} \tablenotetext{e}{Velocity separation in $\mbox{km s}^{-1}$, based MOSFIRE systemic redshifts of the primary and serendip galaxies.} \tablenotetext{f}{Stellar mass ratio between more passive and less massive pair member, irrespective of which galaxy is the ``primary" and which is the ``serendip."} \tablenotetext{g}{Although we identify an object in the F160W image at the location corresponding to where our MOSFIRE serendip spectrum was extracted, there is no 3D-HST catalog entry for this galaxy. We list the ID for this serendip as ``$-$9999" and lack a robust F160W magnitude. We measure its coordinates directly from the F160W image, from which we determine the projected separation for this pair.} \tablenotetext{h}{Primary galaxy with two serendip companions. The primary galaxy is listed twice, once for each serendip.} \tablecomments{The three columns immediately to the right of the primary galaxy ID list the properties of the primary galaxy in each pair, while the corresponding columns immediately to the right of the serendip ID refer to the serendip. The final two columns refer to the properties of the pair.} \end{deluxetable*} \begin{deluxetable*}{lrlrcrrr} \tablecolumns{8} \tablewidth{7in} \tablecaption{Physical Properties of MOSDEF Pairs\label{tab:physical}} \tablehead{ \colhead{Field} & \colhead{ID} & \colhead{Primary/} & \colhead{$\log(M_*)$\tablenotemark{a}} & \colhead{$\log(\mbox{SFR(SED)})$\tablenotemark{b}} & \colhead{$\log(\mbox{SFR(H}\alpha))$\tablenotemark{c}} & \multicolumn{2}{c}{$12+\log(\mbox{O/H})$\tablenotemark{d}} \\ \colhead{} & \colhead{} & \colhead{Serendip} & \colhead{$(M_{\odot})$} & \colhead{$(M_{\odot} \mbox{ yr}^{-1})$} & \colhead{$(M_{\odot} \mbox{ yr}^{-1})$} & \colhead{($O3N2$)} & \colhead{($N2$)} } \startdata \multicolumn{8}{c}{$1.37\leq z \leq 1.70$} \\ \hline COSMOS\tablenotemark{e} & 8381 & primary & 10.94 & 0.60 & \nodata & $<$8.74 & $<$8.88 \\ COSMOS & 8490 & serendip & 10.06 & 0.00 & \nodata & \nodata & \nodata \\ COSMOS & 19325 & primary & 9.45 & 0.76 & 0.40 & $<$8.46 & $<$8.52 \\ COSMOS & 19292 & serendip & 9.18 & 0.56 & $>$0.91 & $<$8.61 & $<$9.00 \\ GOODSN & 10044 & primary & 9.79 & 0.83 & $>$0.15 & \nodata & $<$8.89 \\ GOODSN & 10041 & serendip & 8.75 & -0.21 & $>$0.00 & $<$8.59 & $<$8.69 \\ AEGIS & 35056 & primary & 9.31 & 0.68 & $>$1.21 & $<$8.14 & $<$8.29 \\ AEGIS & 35075 & serendip & 9.28 & 0.32 & 1.13 & $<$8.14 & $<$8.32 \\ AEGIS & 39567 & primary & 10.39 & 1.43 & 1.96 & 8.36 & 8.44 \\ AEGIS & 39897 & serendip & 9.50 & 0.83 & 0.70 & 8.33 & 8.46 \\ AEGIS\tablenotemark{e} & 16339 & primary & 10.55 & 2.25 & 2.36 & 8.61 & 8.58 \\ AEGIS\tablenotemark{e} & 16026 & serendip & 10.83 & 1.87 & 1.70 & \nodata & 8.71 \\ AEGIS & 3237 & primary & 10.57 & 0.88 & $>$1.15 & \nodata & 8.50 \\ AEGIS & 3478 & serendip & 9.66 & 1.75 & 0.95 & 8.51 & 8.61 \\ \hline \multicolumn{8}{c}{$1.90\leq z \leq 2.61$} \\ \hline COSMOS & 1740 & primary & 9.91 & 1.05 & 1.40 & 8.36 & 8.50 \\ COSMOS & 1795 & serendip & 9.27 & -0.25 & $<$1.20 & \nodata & \nodata \\ COSMOS & 10719 & primary & 11.04 & 0.88 & \nodata & \nodata & \nodata \\ COSMOS & 10766 & serendip & 10.21 & 0.69 & $>$0.67 & \nodata & $<$8.78 \\ COSMOS\tablenotemark{e} & 471 & primary & 9.61 & 0.94 & \nodata & \nodata & 8.33 \\ COSMOS & 532 & serendip & 9.03 & 1.55 & \nodata & \nodata & $<$8.60 \\ COSMOS\tablenotemark{e} & 21492 & primary & 11.06 & 2.39 & 2.21 & 8.39 & 8.63 \\ COSMOS & 21598 & serendip & 10.17 & 0.94 & $>$0.37 & $<$8.52 & $<$8.71 \\ COSMOS & 14849 & primary & 10.35 & 1.92 & \nodata & \nodata & $<$8.76 \\ COSMOS\tablenotemark{e} & 14968 & serendip & 9.80 & 1.97 & \nodata & \nodata & 8.70 \\ COSMOS & 25229 & primary & 9.98 & 1.31 & 1.42 & 8.23 & 8.31 \\ COSMOS & $-$9999 & serendip & \nodata & \nodata & \nodata & \nodata & \nodata \\ COSMOS & 7433 & primary & 9.70 & 0.84 & $>$0.83 & $<$8.47 & $<$8.55 \\ COSMOS & 7417 & serendip & 9.40 & 0.79 & $>$0.20 & $<$8.53 & $<$8.67 \\ GOODSN & 17748 & primary & 9.56 & 0.60 & 0.94 & $<$8.25 & $<$8.33 \\ GOODSN & 17714 & serendip & 9.39 & 0.72 & 2.07 & $<$8.07 & $<$8.17 \\ GOODSN & 24825 & primary & 9.75 & 0.43 & \nodata & \nodata & \nodata \\ GOODSN\tablenotemark{e} & 25017 & serendip & 9.65 & 0.70 & 1.36 & \nodata & \nodata \\ GOODSN & 23344 & primary & 10.65 & 1.78 & 2.34 & 8.29 & 8.43 \\ GOODSN & 23418 & serendip & 7.91 & 0.66 & $<$1.37 & \nodata & \nodata \\ GOODSN & 23344 & primary & 10.65 & 1.78 & 2.34 & 8.29 & 8.43 \\ GOODSN & 23339 & serendip & 9.77 & 1.33 & $<$0.95 & \nodata & \nodata \\ GOODSN & 23869 & primary & 10.20 & 1.53 & 1.99 & 8.32 & 8.47 \\ GOODSN & 24074 & serendip & 9.50 & 0.35 & $>$0.45 & $<$8.41 & $<$8.59 \\ GOODSN\tablenotemark{e} & 12302 & primary & 10.92 & 2.11 & 1.99 & 8.56 & 8.68 \\ GOODSN\tablenotemark{e} & 12172 & serendip & 9.69 & 0.32 & 1.40 & 8.36 & 8.45 \\ AEGIS & 18543 & primary & 9.62 & 0.99 & 1.33 & 8.13 & 8.20 \\ AEGIS & 18454 & serendip & 9.33 & 0.37 & $>$1.99 & $<$8.35 & $<$8.58 \\ AEGIS & 31108 & primary & 10.73 & 1.91 & 2.17 & 8.36 & 8.47 \\ AEGIS & 31317 & serendip & 10.00 & 0.68 & \nodata & \nodata & \nodata \\ AEGIS & 36050 & primary & 9.92 & 1.30 & 1.37 & 8.31 & 8.39 \\ AEGIS & 36180 & serendip & 10.51 & 1.72 & 1.47 & 8.51 & 8.60 \\ AEGIS & 29114 & primary & 9.32 & 0.36 & 0.62 & $<$8.26 & $<$8.35 \\ AEGIS & 29045 & serendip & 8.95 & 0.51 & 0.24 & $<$8.57 & $<$8.89 \\ \hline \multicolumn{8}{c}{$2.95\leq z \leq 3.80$} \\ \hline COSMOS & 23183 & primary & 10.06 & 1.82 & \nodata & \nodata & \nodata \\ COSMOS & 23192 & serendip & 10.27 & 1.60 & \nodata & \nodata & \nodata \\ COSMOS & 24579 & primary & 9.42 & 0.80 & \nodata & \nodata & \nodata \\ COSMOS & 24596 & serendip & 8.78 & 0.54 & \nodata & \nodata & \nodata \\ COSMOS & 2360 & primary & 9.73 & 1.06 & \nodata & \nodata & \nodata \\ COSMOS & 2344 & serendip & 9.39 & 0.90 & \nodata & \nodata & \nodata \\ GOODSN\tablenotemark{e} & 28202 & primary & 10.17 & 1.50 & \nodata & \nodata & \nodata \\ GOODSN & 28209 & serendip & 10.12 & 1.45 & \nodata & \nodata & \nodata \\ GOODSN & 15694 & primary & 9.81 & 1.14 & \nodata & \nodata & \nodata \\ GOODSN & 15566 & serendip & 8.83 & 0.34 & \nodata & \nodata & \nodata \\ AEGIS & 30847 & primary & 9.30 & 2.04 & \nodata & \nodata & \nodata \\ AEGIS & 30691 & serendip & 8.81 & 1.56 & \nodata & \nodata & \nodata \enddata \tablenotetext{a}{Log stellar mass, in units of $M_{\odot}$.} \tablenotetext{b}{Log of SFR derived from the best-fit population synthesis model, corrected for dust extinction, in units of $M_{\odot} \mbox{ yr}^{-1}$.} \tablenotetext{c}{Log of SFR derived from H$\alpha$ luminosity, corrected for dust extinction, in units of $M_{\odot} \mbox{ yr}^{-1}$. An entry of ``\nodata" indicates galaxies lacking coverage of either H$\alpha$ or H$\beta$, or with limits in both H$\alpha$ and H$\beta$. SFRs for galaxies with H$\alpha$ detections and H$\beta$ non-detections are indicated as lower limits, while those for galaxies with H$\alpha$ non-detections and H$\beta$ detections are shown with upper limits.} \tablenotetext{d}{Gas-phase oxygen abundance. The column labeled $O3N2$ lists oxygen abundances based on the $O3N2$ indicator, while the column labeled $N2$ contains oxygen abundances based on the $N2$ indicator \citep{pettinipagel2004}. A value of ``\nodata" indicates galaxies lacking coverage of at least one of the required emission lines, or with enough non-detections to prevent derivation of a meaningful limit. Accordingly, we do not report oxygen abundances for galaxies in the highest-redshift bin.} \tablenotetext{e}{Galaxy identified as an AGN based on X-ray, infrared, or rest-optical emission-line properties. Stellar mass, SFR, and metallicity values are not plotted or included in differential comparisons.} \end{deluxetable*} | 18 | 8 | 1808.09978 |
1808 | 1808.03294_arXiv.txt | Extragalactic astronomy relies on the accurate estimation of source photometry corrected for Milky Way dust extinction. This has motivated the creation of a number of ``Galactic" dust maps. We investigate whether these maps are contaminated by extragalactic signals using the clustering-redshift technique, i.e., by measuring a set of angular cross-correlations with spectroscopic objects as a function of redshift. Our tomographic analysis reveals imprints of extragalactic large-scale structure patterns in nine out of 10 Galactic dust maps, including all infrared-based maps as well as ``stellar'' reddening maps. When such maps are used for extinction corrections, this extragalactic contamination introduces redshift- and scale-dependent biases in photometric estimates at the millimagnitude level. It can affect both object-based analyses, such as the estimation of the Hubble diagram with supernovae, as well as spatial statistics. The bias can be appreciable when measuring angular correlation functions with low amplitudes, such as lensing-induced correlations or angular correlations for sources distributed over a broad redshift range. As expected, we do not detect any extragalactic contamination for the dust map inferred from 21cm HI observations. Such a map provides an alternative to widely used infrared-based maps but relies on the assumption of a constant dust-to-gas ratio. We note that, using the WISE 12 micron map sensitive to polycyclic aromatic hydrocarbons (PAH), an indirect dust tracer, we detect the diffuse extragalactic PAH background up to $z\sim2$. Finally, we provide a procedure to minimize the level of biased magnitude corrections in maps with extragalactic imprints. | \label{sec:intro} \begin{figure*}[t] \centering \includegraphics[width=0.71\textwidth]{f01.pdf} \caption{ Bottom panel: spectral features of dust extinction \citep[dashed line; $A_{\lambda}$ from][]{2001ApJ...548..296W} and emission \citep[solid line; ${\rm log}\ j_{\nu}$ from][]{2007ApJ...657..810D}, together with a delta function for HI 21cm emission. The arrow shows the expected shift when emission temperature or redshift changes. A set of filter curves from several surveys is shown above with arbitrary normalization. Top panel: summary of spectral sampling for the 10 Galactic dust-reddening maps considered in this paper. } \label{fig:dust_observables} \end{figure*} Our view of the extragalactic sky is altered by the presence of dust in the interstellar medium (ISM) of the Milky Way. Dust grains absorb and scatter incident photons at short wavelengths and emit radiation at longer wavelengths. Collectively, they produce a foreground screen that extincts and reddens the light of extragalactic objects. Correcting for such extinction effects is important for the accurate estimation of extragalactic source photometry, and the need for higher accuracy keeps increasing from the demanding requirements of precision cosmological experiments. Dust can be traced over a wide range of wavelengths. This is illustrated in Figure~\ref{fig:dust_observables} where we show the expected extinction by Milky Way-type dust in the ultraviolet to near-infrared (IR) from \cite{2001ApJ...548..296W}, and the dust thermal emission from the model of \cite{2007ApJ...657..810D} with a broad peak in the far-IR and polycyclic aromatic hydrocarbon (PAH) features in the mid-IR. As dust is usually mixed with neutral gas, we also show the 21cm hydrogen emission line, which has been used as an indirect dust tracer. This latter technique was used to introduce the first large-scale dust map by \cite{1978ApJ...225...40B,1982AJ.....87.1165B}. Later, the $IRAS$ and $COBE$ satellites launched in the 1980s opened up the atmospherically opaque window in the mid- to far-IR. \citet[][hereafter SFD]{1998ApJ...500..525S} used this data to create a dust-reddening map based on the thermal continuum emission of dust. Additional dust maps have been derived using IR data from the $Planck$ satellite, PAH emission using the Wide-field Infrared Survey Explorer ($WISE$), and optical reddening toward background stars \citep{2012ApJ...757..166B} or galaxies \citep{2010ApJ...719..415P}. The tight correlation between dust reddening and some of the diffuse interstellar bands has also been demonstrated and used to map Galactic dust in the optical \citep{2015MNRAS.452.3629L} and the near-IR \citep{2015ApJ...798...35Z}. Due to the broad width of the blackbody spectrum, Galactic dust maps derived from IR emission measurements unavoidably include source contributions from extragalactic objects over a wide range of redshifts. Such maps therefore suffer from some level of extragalactic contamination. Such a signature has been reported by \cite{2007PASJ...59..205Y} who analyzed the number counts of low-redshift galaxies from the Sloan Digital Sky Survey (SDSS) as a function of the Galactic reddening measured in SFD. They showed that, toward low-reddening sight lines, the reddening value $E_{B-V}$ in SFD appears to be roughly proportional to the number density of galaxies, whereas it should not depend on such a quantity. Motivated by this result, \cite{2013PASJ...65...43K} showed that the mean reddening values derived by SFD at the locations of SDSS photometric galaxies and quasars are in excess with respect to their angular vicinity. This excess reddening can be largely accounted for by the contamination in far-IR emission from both the stacked galaxies and the galaxies clustered around them. \citet{2015MNRAS.446.2696S} applied the clustering-based redshift estimation technique to probe the extragalactic contribution of various intensity maps produced by the Planck Collaboration. These authors generalized the clustering-redshift technique (originally introduced for discrete sources by \citealt{2008ApJ...684...88N} and \citealt{2013arXiv1303.4722M}) to diffuse fields. This allowed them to reveal extragalactic contributions in the $Planck$ dust map. In this paper, we follow a similar line of investigation. We analyze 10 Galactic dust maps available in the literature and perform a systematic clustering-redshift analysis for each of them in a uniform manner. We constrain the angular and redshift dependence of the extragalactic imprints in each dust map by measuring a set of angular cross-correlations between these maps with a spectroscopic reference sample of galaxies and quasars as a function of redshift. In Section~\ref{sec:formalism} we lay out the formalism of extinction correction, the impact of a biased reddening associated with extragalactic imprints on source number counts and clustering statistics, and how to measure it using cross-correlations. In Section~\ref{sec:data} we introduce the dust maps and cross-correlation reference sample. In Section~\ref{sec:exgalresults}, we present the results in measuring the amplitudes and redshift dependence of extragalactic imprints in each dust map. We discuss the implications for precision cosmology and summarize our results in Section~\ref{sec:discussion}, and \ref{sec:summary}, respectively. A flat universe with Planck cosmological parameters is assumed \citep{2014A&A...571A..16P}. | \label{sec:number_counts_and_clustering_biases} Let ${N}(\phi,\,m)$ be the differential number count of a population of extragalactic sources in the magnitude range [$m$, $m+dm$] and angular space [$\phi$, $\phi+d\Omega$]. If magnitude estimates are biased with a shift $\delta m(\phi)$, the apparent number count reads \begin{equation} {\hat{N}_{\rm {dered}}}(m) = {\rm N}(m+ \delta m )\;. \label{eq:number_counts_a} \end{equation} Considering a power-law luminosity function with a slope \begin{equation} \alpha(m)=2.5\, \frac{ \textrm{d\,log}\, {N}(m)}{{\rm d}m}\,, \label{eq:alpha} \end{equation} and expressing the magnitude shift in terms of a flux ratio $\mu (\phi)$ such that \begin{equation} \mu = 10^{-0.4\,\delta m} \approx 1-0.92\, \delta m \;, \label{eq:mu} \end{equation} Equation~\ref{eq:number_counts_a} becomes \begin{eqnarray} \hat{N}_{\rm dered}(m)\, &=&\, \mu^{\alpha}\, {N}(m) \nonumber\\ &\approx&\, (1+\alpha\, \delta \mu)\, {N}(m)\;, \label{eq:number_counts} \end{eqnarray} where the second equality is based on $\mu = 1+\delta \mu$, with fluctuations $\delta\mu$ being small compared to unity. Note that the slope of source number counts $\alpha$ is typically of order unity. To summarize, we have \begin{equation} - \delta m \approx \delta \mu \approx R_{\lambda}\, \langle \delta E_{B-V}\rangle_g \approx \delta A^{\rm EG}_{\lambda} \;. \end{equation} We point out that this formalism describing changes in number counts under brightness changes is similar to that used in gravitational-lensing magnification. Both extinction overcorrection and gravitational magnification introduce biases in number counts, density field, and angular clustering measurements for magnitude-limited samples. Considering source overdensity fluctuations over the sky, \begin{equation} \delta(\phi) = {N}(\phi)/\langle {N}\rangle -1\; , \label{eq:overdensity_def} \end{equation} Equation~\ref{eq:number_counts} shows that a biased dereddening estimate will introduce modulations in the apparent number counts. The estimated overdensity is given by \begin{eqnarray} \hat{\delta}_{{\rm dered}}(\phi) \, &=& \, \delta(\phi)[1+\alpha\, \delta \mu(\phi)]+\alpha\, \delta \mu(\phi)\nonumber\\ \,&\simeq&\, \delta(\phi)+\alpha\, \delta \mu(\phi)\,, \label{eq:apparent_delta_g} \end{eqnarray} where the $\delta \mu$ modulation could be appreciable when the density contrast $\delta$ is small. Any magnitude-limited sample of dereddened sources will thus carry spatial fluctuations of the extragalactic imprints present in the dust map used for extinction correction. As a direct consequence, angular correlations between two magnitude-limited, dereddened populations will also be affected. Let us consider the density fields $\delta_1$ and $\delta_2$ for these two populations with an intrinsic angular clustering $w_{12}= \langle \delta_{1} \cdot \delta_{2}\rangle$. Following Equation~\ref{eq:apparent_delta_g}, the bias field in the dust map will modulate the apparent two-point function such that \begin{eqnarray} \hat{w}_{12, \rm dered} \, &=&\, \langle \hat{\delta}_{1, \rm dered} \cdot \hat{\delta}_{2, \rm dered}\rangle \nonumber\\ &\simeq&\, w_{12} + \alpha_2\, \langle \delta_1 \cdot \delta\mu \rangle + \alpha_1\, \langle \delta_2 \cdot \delta\mu \rangle \nonumber\\ &+& \alpha_1\,\alpha_2\, \langle \delta\mu^2 \rangle\;, \label{eq:2pt_fn_bias} \end{eqnarray} where the second and third terms are the source--dust map bias correlations (one for each population), and the last term is the autocorrelation of the extragalactic bias in a dust map. The presence of extragalactic contamination thus sets a floor affecting two-point function measurements using magnitude-limited, dereddened extragalactic sources. \subsection{Measuring Extragalatic Imprints in a Dust Map} We now consider certain properties of dust maps relevant to our analysis. \begin{enumerate} \item Dust maps are typically provided using reddening units; hereafter, we will follow this convention and express dust column density estimations in units of $E_{B-V}$. \item If the zero point of the reddening is properly calibrated in a map, the spatial average of the extragalactic imprints (e.g., the cosmic IR background (CIB) monopole in IR-based maps) should already be removed. The information that can be extracted is therefore in the fluctuations of the reddening field, \begin{eqnarray} \delta E_{B-V}(\phi) = \delta E_{B-V}^G(\phi)+ \delta E_{B-V}^{EG}(\phi)\nonumber \\ \equiv E_{B-V}(\phi) - \langle E_{B-V} (\phi)\rangle, \label{eq:delta_ebv} \end{eqnarray} where $\langle E_{B-V}\rangle$ is the sky average of the reddening field. As our goal is to extract extragalactic signals, $\delta E_{B-V}^{EG}$, we choose to estimate $\langle E_{B-V} \rangle$ using a running mean with a radius of $1^{\circ}$ to suppress Galactic contributions. \item The angular autocorrelation of dust maps $\langle (\delta E_{B-V})^2 (\theta) \rangle$ is often of limited interest, as it mixes both Galactic and extragalactic fluctuations. This autocorrelation can only be used as an upper limit to the level of extragalactic contamination. \item It is possible to separate the extragalactic component from the Galactic reddening by considering angular cross-correlations between the dust map and an external set of reference galaxies. Let $\delta_r (\phi, z)$ be the fractional overdensity field of these reference galaxies. Since the spatial distribution of dust in the Milky Way Galaxy is not expected to correlate with such a population, the reference--dust-reddening correlation reads \begin{eqnarray} \langle \delta_r(\phi, z) \cdot \delta E_{B-V}(\phi+\theta) \rangle\, &=&\, \langle \delta E_{B-V}^{\rm EG}(\theta, z) \rangle_r\, , \label{eq:cross-correlation} \end{eqnarray} where the subscript $r$ refers to the positions of the reference objects. This excess-reddening estimator thus allows us to extract extragalactic imprints as a function of redshift in each of the dust maps considered. To enhance the signal-to-noise ratio of this estimator, we also construct a variance map for each $\delta E_{B-V}$ field and replace the mean with an inverse variance-weighted mean for the $\langle...\rangle$ operator (see Appendix A for details). \end{enumerate} \label{sec:discussion} We have shown that, out of the ten wide-field Galactic dust maps currently available, nine present detectable extragalactic contamination. The remaining one, based on the hydrogen distribution, relies on an assumed dust-to-gas ratio that can be shown to have a complex spatial distribution over the sky. What are the impacts of these extinction overcorrection biases on astronomical experiments? As presented in section~\ref{sec:number_counts_and_clustering_biases}, a reddening overcorrection leads to a whole hierarchy of number count bias, overdensity bias, and clustering bias (Equations~\ref{eq:number_counts}--\ref{eq:2pt_fn_bias}). The leading terms of these biases are all of the same order: $R_{\lambda}\, \langle \delta E_{B-V}\rangle_g \approx -\delta m \approx \alpha\,\delta \mu$. The exact amplitude of this effect depends on the chosen dust map, wavelength, redshift, source population, and angular scale considered. In many cases, they are found to be at the millimagnitude level, i.e. of order $10^{-3}$ in the optical, assuming $R_{\lambda} \approx 3$ and $\alpha$ of order unity. \begin{figure*}[!th] \includegraphics[width=0.98\textwidth]{f09.pdf} \caption{Case comparison for a galaxy--quasar clustering introduced by gravitational magnification (black curves; theoretical expectations) and extinction overcorrection (red curves for the $V$-band SFD bias scaled from our $\Delta E_{B-V}$ measurements). The redshifts of the galaxies and quasars are set to 0.2 and 1.5, respectively. The results are shown for a range of $\alpha_q$, the logarithmic slope of the quasar magnitude distribution, where $\alpha_q-1=0.2$ (second panel) is roughly the effective value for an optimal estimator using quasars with $g<21$.} \label{fig:lensing} \end{figure*} Systematic shifts of order millimagnitude can potentially impact precision cosmology experiments, for example, using standard candles, where the sample mean brightness has to be measured with a fractional error comparable to the targeted precision in certain cosmological parameters. Similarly, photometric offsets can affect clustering measurements such as spatial correlation functions, which are used for a wide range of applications including galaxy--halo connection, baryonic acoustic oscillations, gravitational lensing, neutrino masses, cold/warm dark matter, etc., all requiring extinction correction for the tracer sample of the experiments (various kinds of galaxies), while not all of them will be affected by a small bias. Equation~\ref{eq:apparent_delta_g} provides a rule of thumb to identify the regime in which extinction overcorrection bias will be significant: since the bias in the source overdensity $\delta (\phi)$ is roughly $\alpha\, \delta \mu(\phi)$, the bias is only important in a weak-field regime where $\delta (\phi)$ is not much greater than $\alpha\, \delta \mu(\phi)\sim 10^{-3}$ (in the optical). For a typical galaxy overdensity in 3D, $\delta (\phi)$ fluctuates at a level greater than unity on arcminute scales. However, this amplitude decreases with an increasing level of line-of-sight projection; experiments involving a substantial spread in redshift, either imposed by the photometric selection or due to photometric redshift errors, will end up falling in that category. Another relevant context is gravitational lensing-induced correlation functions, which we discuss in more detail below. Finally, we can also imagine a limiting case with a hypothetical population of randomly distributed sources (i.e., zero intrinsic clustering). When estimating the correlation function of such a population after correcting for Galactic extinction, one will end up measuring the clustering of the extragalactic contamination imprinted in the dust map. \begin{figure*}[t!] \includegraphics[width=0.98\textwidth]{f10.pdf}\\ \caption{Changes in distance modulus after correcting for a rest-frame $B$-band extinction overcorrection in SFD (red data points). Black lines overlay the changes corresponding to perturbations in each of the cosmological parameters $H_0$ (left), $\Omega_m$ (middle), and $w$ (right) from a fiducial cosmology (blue dash-dotted lines). After correcting for the extinction overcorrection, one expects $H_0$ to increase by $<0.1\%$ and $\Omega_m$ and $w$ to decrease by a fraction of a percent.} \label{fig:sn} \end{figure*} \subsection{Impact on Weak-field Clustering: Lensing Magnification}\label{sec:discussion_lensing} Gravitational-lensing magnification modulates the apparent number count of background sources by two mechanisms: magnitude brightening and solid-angle dilatation. Analogs to Equation~\ref{eq:number_counts}, this is expressed by \begin{equation} {\rm N}_{\rm obs}(m)= [1+(\alpha - 1)\, \delta \mu]\, {\rm N}(m) \label{eq:number_counts_lensing} \end{equation} \citep[e.g.,][]{2005ApJ...633..589S}, where the slope of the source number counts $\alpha$ is defined in Equation~\ref{eq:alpha}, the $-1$ term takes into account the source dilution due to area dilatation, and $\delta \mu$ is the lensing magnification factor. This is similar to the formalism introduced in Section~\ref{sec:formalism} for the extinction overcorrection. The only difference is the change in sky solid angle. The magnification effect leads to apparent angular correlations between foreground lenses and background sources that are physically uncorrelated. This can be expressed as a special case of Equation~\ref{eq:2pt_fn_bias}, \begin{eqnarray} w_{fb, \rm obs}(\theta) \, &=&\, \langle \delta_{f}(\phi) \cdot \delta_{b, \rm obs}(\phi+\theta)\rangle \\ &=&\, (\alpha_b-1)\, \langle \delta_f (\phi) \cdot \delta \mu (\phi+\theta)\rangle \nonumber\\ &=&\, (\alpha_b-1)\, b_f\, w_{\mu m}(\theta)\,, \label{eq:2pt_fn_bias-lensing} \end{eqnarray} where only the background overdensity is modulated, and the third equality assumes a linear bias $b_f$ relating $\delta_f$ to matter overdensity $\delta_m$. The gravitational potential associated with the foreground lenses introduces a magnification field $w_{\mu m}(\theta)$ to the sources, whose expression can be found in, e.g., \cite{2001PhR...340..291B}. Depending on the value of $\alpha_b - 1$, lensing-induced clustering can be positive, zero, or negative. Table~\ref{table:lensing} summarizes the effects of extinction overcorrection and lensing magnification in parallel, and one can see the similarities between the two. Unlike lensing, however, extinction overcorrection is chromatic. Another difference between the two is that the lensing efficiency is a function of the angular diameter distances of both populations, while the line-of-sight efficiency of extinction overcorrection is a constant as one applies such a correction using 2D dust maps. Given this constant efficiency, extinction overcorrection also affects source autocorrelations. The extinction overcorrection will affect measurements of lensing-induced correlation functions. To estimate the level at which these effects occur, we examine the following scenario commonly targeted in lensing studies (see, for example, \citealt{2005ApJ...633..589S}). We consider a population of foreground galaxies at $z=0.2$ and background quasars at $z=1.5$. The magnification-induced angular correlation function expected between these two populations is shown in Figure~\ref{fig:lensing} (black curves) for several magnitude bins that correspond to several values of $\alpha_q-1$ following Equation~\ref{eq:2pt_fn_bias-lensing}. Again, $\alpha_q$ is given by the shape of the quasar luminosity function following Equation~\ref{eq:alpha}, and the range of $\alpha_q-1$ in Figure~\ref{fig:lensing} corresponds to that for quasars from $g$ of 17--21 mag going from the left to right panels. Red curves show the expected $w_{gq}$ introduced by extinction or dereddening overcorrection in the $V$ band if the SFD map is used. This is based on our excess-reddening measurements presented in Section~\ref{sec:exgalresults} assuming an $R_V$ of 3.1. As provided in Equation~\ref{eq:2pt_fn_bias}, the extinction correction--induced correlation has two terms: one that scales with $\alpha_q$ as in lensing and an additional $\alpha_g$ term. Here we set $\alpha_g$ to 1 for all panels. For large $\alpha_q$ (left panels), lensing dominates over extinction overcorrection. When $\alpha_q$ approaches 1, however, lensing effects vanish, while the extinction overcorrection is only slightly reduced, since the $\alpha_g$ term has not changed. The extinction overcorrection can thus bias lensing magnification measurements, especially at the faint end (right panels). For an optimal lensing estimator weighted by the number of quasars and the expected signal using quasars brighter than $g$ of $21$ (SDSS depth), the effective $\alpha_q-1$ is about 0.2 \citep{2002A&A...386..784M,2005ApJ...633..589S}. This is shown in the second panel, where the extinction correction--induced correlation is about $50\%$ of that induced by lensing. We thus already expect some impacts biasing current lensing measurements. Upcoming surveys such as LSST \citep{2009arXiv0912.0201L}, $Euclid$ \citep{2011arXiv1110.3193L} and $WFIRST$ \citep{2013arXiv1305.5422S} will provide us with photometric samples enabling lensing-induced correlation measurements with a precision largely surpassing that of existing measurements. In this regime, the extinction overcorrection discussed above will become a significant limitation in harnessing the full statistical power of the expected datasets. It will be important to correct or take this effect into account. \begin{figure*}[t!] \begin{center} \includegraphics[width=0.68\textwidth]{f11.pdf} \end{center} \caption{Normalized excess $E_{B-V}$ (within $10'$) for all 10 reddening maps. The effect of our choice of reference samples has been removed by dividing out the linear galaxy bias. One can use these values for generic corrections of the dust extinction overcorrection bias present in each map.} \label{fig:normalized_delta_ebv} \end{figure*} \subsection{Cosmological Parameters with Standard Candles}\label{sec:discussionSNIa} The extinction overcorrection is expected to impact cosmological parameter extractions using Type Ia supernovae as standard candles. Such studies utilize the cosmology dependence of the luminosity distance, thus distance modulus as a function of redshift \citep[see a review in][]{2011ARNPS..61..251G}. Briefly, the distance modulus\footnote{Unfortunately, by convention, the notation $\mu$ is the same with that for the magnification factor in the context of lensing.} $\mu$ of a standard candle with an absolute magnitude $M$ can be used to probe its luminosity distance $D_{\rm L}$: \begin{equation} \mu = m - M =5\, \textrm{log} \left(\frac{D_{\rm L}(z)}{10~\textrm{pc}}\right), \end{equation} where $m$ is its apparent magnitude. The luminosity distance $D_{\rm L}$ as a function of redshift is cosmology-dependent, \begin{equation} D_{\rm L}(z) = \frac{c}{H_0}\, (1+z)\, \int_0^z\frac{dz'}{E(z')}, \end{equation} where $H_0$ is the Hubble constant and \begin{equation} E(z)=\sqrt{\Omega_{m}\, (1+z)^{3} + \Omega_{\Lambda}\, (1+z)^{3\, (1+w)}} \end{equation} is the normalized Hubble function for a flat $\rm \Lambda CDM$ universe. In this expression, $\Omega_{m}$ and $\Omega_{\Lambda}$ are the present-day matter and dark energy density in the unit of the critical density, respectively, and $w$ is the dark energy equation-of-state parameter. As the distance modulus $\mu$ is estimated after correcting for Galactic extinction, one expects an extinction overcorrection bias $\delta m(\phi) = - R_{\lambda}\, \delta E_{B-V}^{\rm EG}(\phi)$ (Equation~\ref{eq:ext_overcorrection}). This bias is redshift-dependent and can therefore impact cosmological parameter estimation. The estimated supernova dereddened distance modulus is, on average, \begin{eqnarray} \langle \hat{\mu} \rangle_{\rm SN, dered}\, &=&\, \langle \mu \rangle_{\rm true} + \langle \delta m \rangle_{\rm SN}(z) \nonumber\\ \, &=&\, 5\, \textrm{log} \left(\frac{\langle D_{\rm L}\rangle (z)}{10~\textrm{pc}}\right) - R_{\lambda}\, \langle \delta E_{B-V}^{\rm EG}\rangle_{\rm SN}(z)\,, \end{eqnarray} where $\langle \delta E_{B-V}^{\rm EG}\rangle_{\rm SN}$, the excess reddening due to extragalactic imprints in the dust map, is to be evaluated at the zero lag ($\theta = 0$) toward the supernovae. Here we investigate the fractional biases in $H_0$, $\Omega_m$, and $w$ under the level of extinction overcorrection effect measured in Section~\ref{sec:exgalresults}. For simplicity, we assume that the supernova hosts are galaxies similar to our SDSS reference objects, i.e., $\langle ... \rangle_{\rm SN} = \langle ... \rangle_{r}$. Based on the angular dependence of $ \langle \delta E_{B-V} \rangle_r$ we found that, as shown in Figure~\ref{fig:angular-correlation}, for maps with about 5 arcmin resolution like SFD and the $Planck$ maps, the zero lag $ \langle \delta E_{B-V} \rangle_r (\theta=0)$ is roughly 2 times the $10'$ averaged $\Delta E_{B-V}$ presented in Figure~\ref{fig:excess-ebv-redshift}. For the wavelength dependence $R_{\lambda}$, we adopt an extinction vector calibrated in \cite{2011ApJ...737..103S} assuming an $R_{V}$ of $3.1$ (this extinction vector includes a correction for a 13\% overestimation of the $E_{B-V}$ in SFD). Figure~\ref{fig:sn} shows the distance modulus offset $\delta \mu$ as a function of redshift after we correct for an extinction overcorrection for the SFD map in the rest-frame $B$ band (red data points). As is typically done in supernova cosmology studies, since there is an uncertainty in calibrating the distance ladder, we anchor the $\delta \mu$ to zero locally at our first redshift bin, $z\sim0.02$. Over a redshift range of order unity, extinction overcorrection thus changes the standard candles by 1.5 mmag. Figure~\ref{fig:sn} also overplots the distance modulus offset once we perturb each of the three cosmological parameters in each panel with the labeled fractional changes (black lines) from a fiducial flat universe cosmology of $[H_0,\ \Omega_m,\ w]=[70, 0.3, -1]$ (blue dash-dotted lines). We find that the bias given by the mean $\delta m$ in $H_0$ is small (less than $0.1\%$). For $\Omega_m$ and $w$ where a $0.5\%$ bias is expected, the effect in extinction overcorrection is going to be important when upcoming cosmology experiments are targeting $1\%$ level precision. Among these three cosmological parameters, accurate measurements of $w$ are of critical importance in the coming decade, as a significant departure of $w$ from unity would rule out the scenario of dark energy being a cosmological constant. We therefore suggest incorporating the correction of extinction-correction bias ($\sim0.5\%$) in upcoming Hubble diagram estimates. \subsection{Bias Correction}\label{sec:correct_bias} We now describe a procedure to correct for the biased dereddened magnitude estimations due to extragalactic imprints in Galactic dust maps. For a population of objects at a given redshift, the $10'$ average reddening excess $\Delta E_{B-V}$ presented in Figure~\ref{fig:excess-ebv-redshift} provides a starting point to quantify the amount of extinction overcorrection when a given dust map is used. Since the $\Delta E_{B-V}$ in Figure~\ref{fig:excess-ebv-redshift} is measured around specific types of reference objects (SDSS spectroscopic galaxies and quasars), for other types of objects, the mean excess reddening needs to be scaled with the clustering bias ratio. In Figure~\ref{fig:normalized_delta_ebv} we provide $\Delta E_{B-V}/b_r$, the excess-reddening estimations as a function of redshift for each map, similar to that present in Section~\ref{sec:exgalresults} but this time normalized by the bias of our reference sample (see Appendix B for the $b_r$ measurements). With this information, we can determine the mean magnitude bias of an arbitrary galaxy population $g$ at redshift $z$ with a linear clustering bias $b_g$ following Equation~\ref{eq:pop_mag_shift}: \begin{eqnarray} \langle \delta m (z) \rangle_g \, &=&\, - R_{\lambda}\, \langle \delta E_{B-V} (z, \theta = 0) \rangle_g \nonumber\\ &=&\, -R_{\lambda}\, C\, \frac{\Delta E_{B-V}(z)}{b_r}\, b_g, \end{eqnarray} where we make clear that, since magnitude is a one-point statistic, the offset is related to the excess reddening at zero lag ($\theta = 0$). The constant $C$ is a beam correction to relate the $\Delta E_{B-V}$ measured within $10'$ to the zero-lag reddening excess. For maps like SFD or $Planck$ of about $5'$ resolution, $C\approx 2$, which can be visualized in Figure~\ref{fig:angular-correlation}; for lower-resolution maps with a spatial half width at half maximum compatible or larger than $10'$, $C\approx 1$ (thus, no beam correction is needed). For convenience, we provide a fitting function for the magnitude bias in the SFD map, \begin{eqnarray} \langle \delta m^{\textrm{SFD}}(z)\rangle_g\, &=&\, -R_{\lambda}\, \langle \delta E^{\textrm{SFD}}_{B-V}(z, \theta=0) \rangle_g \nonumber \\ &\simeq&\, -0.024\, R_{\lambda}\,b_g\,(z+0.16)^{-1.8}\;, \end{eqnarray} in mmag. To correct for this extinction overcorrection, one simply subtracts the $\langle \delta m \rangle$ (adds a positive magnitude) from the estimated dereddened magnitudes. We have analyzed 10 Galactic dust maps and investigated whether they are contaminated by extragalactic signals. Our tomographic analysis, based on the so-called clustering-redshift technique, has shown that 9/10 dust maps present imprints of extragalactic large-scale structure patterns, in some cases detected up to $z\sim4$. These extragalactic signals are found in all IR-based maps, from $12$ $\mu$m to the millimeter range, as well as ``stellar'' optical reddening maps. The amplitude of this extragalactic contamination is typically found to be at the millimagnitude level. Its redshift and angular scale variations depend on the chosen dust map. More specifically, we find the following. \begin{enumerate} \item for all IR thermal dust maps, including the widely used \cite{1998ApJ...500..525S} map, Galactic reddening $E_{B-V}$ is systematically over-estimated around galaxies and quasars up to $z\sim2$ at a level ranging from subpercent to a few percent on scales of $10'$. This originates from CIB fluctuations due to the emission from dusty star-forming galaxies. The more recent \cite{2014A&A...571A..11P, 2016A&A...586A.132P} dust maps present a higher level of extragalactic contamination at higher redshift as they probe the Rayleigh-Jeans side of the dust blackbody emission spectrum, resulting in a negative $K$-correction. In addition, at $z>2$, this effect is further enhanced due to the peak of the cosmic star-formation history. \item For the stellar reddening maps using point-source optical photometry in Pan-STARRS1, we find an underestimation of Galactic reddening, especially around quasars at $1<z<1.5$ at the percent level. This reveals issues in star--galaxy and/or star--quasar separations. \item The $WISE$ 12$\mu$m map from \cite{2014ApJ...781....5M} is sensitive to PAH emission and has been used to create a Galactic dust map based on this tracer. Analyzing it, we detect the diffuse extragalactic PAH background up to $z\sim2$ and find that the Galactic PAH-to-dust ratio is similar to the cosmic mean. \item We have found the HI-based reddening map from \cite{2017ApJ...846...38L} to be free of extragalactic contamination, at least down to the $5\times10^{-4}$ level. Such a map provides an alternative to the more standard IR-based dust maps but relies on an assumed dust-to-gas ratio, whose spatial fluctuations can lead to an error of about 3 mmag in $E_{B-V}$. \end{enumerate} When these maps are used for correcting the photometry of extragalactic objects for Milky Way extinction, redshift- and scale-dependent biases are introduced. These artificial magnitude offsets then lead to biases in galaxy number counts and spatial auto- and cross-correlations at a level of about $10^{-3}$--$10^{-2}$ on scales of $10'$. These effects can then impact precision cosmology experiments. They can affect both object-based analyses and spatial statistics. These biases can be appreciable when estimating angular correlation functions with low amplitudes such as lensing-induced correlations or angular correlations for sources distributed over a broad redshift range. For precision cosmology with Type Ia supernovae, we expect a $0.5\%$ impact on the determinations of $\Omega_m$ and $w$, which will be significant for upcoming surveys like LSST and WFIRST targeting the one percent precision range. For such experiments, we recommend testing the robustness of the final results against the different dust maps used. Finally, we provide a procedure to correct for or decrease the level of biased magnitude corrections in maps with extragalactic imprints. | 18 | 8 | 1808.03294 |
1808 | 1808.01402_arXiv.txt | We have obtained high-resolution spectra of Jupiter's Great Red Spot (GRS) between 4.6 and 5.4~\um~ using telescopes on Mauna Kea in order to derive gas abundances and to constrain its cloud structure between 0.5 and 5~bars. We used line profiles of deuterated methane (\dm)~ at 4.66~\um~ to infer the presence of an opaque cloud \textbf{at 5$\pm1$~bars}. From thermochemical models this \textbf{is almost certainly} a water cloud. We also used the strength of Fraunhofer lines in the GRS to obtain the ratio of reflected sunlight to thermal emission. The level of the reflecting layer was constrained to be at \textbf{570$\pm30$~mbars} based on fitting strong \ammonia~ lines at 5.32~\um. We identify this layer as an ammonia cloud based on the temperature where gaseous \ammonia~ condenses. We found evidence for a strongly absorbing, but not totally opaque, cloud layer at pressures deeper than 1.3 bars by combining Cassini/CIRS spectra of the GRS at 7.18~\um~ with ground-based spectra at 5~\um. This is consistent with the predicted level of an \amhs~ cloud. We also constrained the vertical profile of \water~ and \ammonia. The GRS spectrum is matched by a saturated \water~ profile above an opaque water cloud at 5~bars. The pressure of the water cloud constrains Jupiter's O/H ratio to be at least 1.1 times solar. The \ammonia~ mole fraction is \textbf{200$\pm50$~ppm} for pressures between 0.7 and 5~bars. Its abundance is \textbf{40~ppm} at the estimated pressure of the reflecting layer. We obtained \textbf{0.8$\pm0.2$~ppm} for \ph, a factor of 2 higher than in the warm collar surrounding the GRS. We detected all 5 naturally occurring isotopes of germanium in \germane~ in the Great Red Spot. We obtained an average value of \textbf{0.35$\pm0.05$~ppb} for \germane. Finally, we measured \textbf{0.8$\pm0.2$~ppb} for CO in the deep atmosphere. | The Great Red Spot is a high-pressure region in the atmosphere of Jupiter, producing an anticyclonic storm 22\deg~ south of the planet's equator. The spot is large enough to contain two or three planets the size of Earth. There have been numerous studies of various aspects of the Great Red Spot (GRS) including its wind field (\citet{simon02}, \citet{asay09}), its mysterious red color (\citet{loeffler16}, \citet{carlson16}), its significant shrinkage in size (\citet{asay09}, \citet{simon14}) and its overall dynamics (e.g. \citet{marcus93}, \citet{read06}, \citet{palotai14}). There have also been numerous studies of the vertical cloud structure of the GRS. \citet{banfield98} used Galileo Solid State Imaging (SSI) data between 0.73 and 0.89~\um~ to infer a thick cloud over the GRS extending from 200 to 700~mbars. \citet{irwin99} reported a similar cloud structure for the GRS based on \citet{baines96}. These authors used Galileo Near Infrared Mapping Spectrometer (NIMS) spectra of the GRS between 0.89 and 2.2~\um. This cloud forms at the level where \ammonia~ ice is expected to condense. However, \citet{baines02} found that spectrally identifiable ammonia clouds are present to the northwest of the GRS, but not in the spot itself. Clouds to the northwest were newly condensed and therefore fresh, while in the GRS cloud particles \textbf{may have evolved} through coating or compositional mixing (e.g. \citet{atreya05, kalogerakis08, sromovsky10}). \citet{simon01, simon02} used Galileo SSI data between 0.41 and 0.89~\um~ to infer the presence of an optically thin stratospheric haze, a moderate to dense tropospheric haze, and an optically thick, physically thin cloud sheet around 900~mbars. Prior studies of gas composition in the GRS have mostly been limited to the upper levels of the feature, between about 100 - 500~mbar, due to opacity from aerosols and \textbf{Rayleigh} scattering. Retrievals of locally-depleted NH$_3$ in the 100 - 300~mbar range from Voyager IRIS \citep{griffith92} and HST/FOS \citep{edgington99} might seem to be at odds with locally-enhanced NH$_3$ abundance near the 500-mbar level from Cassini/CIRS \citep{achterberg06}. These results might be consistent if NH$_3$ decreases rapidly with height (0.5-km scale height), as found by \citet{tokunaga80} based on ground-based infrared spectroscopy. \citet{fletcher10} used multiple sources of thermal infrared data to show that there are also horizontal gradients of composition (as well as temperature and aerosol parameters) across the GRS. \citet{fletcher16} revisited the GRS using ground-based infrared spectral imaging with TEXES, confirming the north-south gradient of NH$_3$ concentration in the 500-mbar region. At radio frequencies, the Very Large Array (VLA) has been used to produce spectral maps of NH$_3$ concentrations at deeper levels on Jupiter, including the GRS \citep{depater16}. Wavelengths from 1.7 - 5.5~cm probe the 0.6 - 6 bar pressure levels, where the GRS has higher NH$_3$ concentrations than most other regions of the planet (except for the Equatorial Zone and plumes of NH$_3$-rich gas near 4$^{\circ}$N). At these wavelengths, gaseous NH$_3$ is the principal opacity source, while cloud opacity is expected to be minimal. In this paper we present ground-based observations of Jupiter's Great Red Spot between 4.6 and 5.4 \um. The 5-\um~ region is a window to the deep atmosphere of Jupiter because of a minimum in opacity due to \h~ and \methane. This spectrum provides a wealth of information about the gas composition and cloud structure of the troposphere. Jupiter's 5-\um~spectrum is a mixture of scattered sunlight and thermal emission that varies significantly between Hot Spots and low-flux regions such as the Great Red Spot. Chemical models of Jupiter's cloud structure predict three distinct layers: an \ammonia~ice cloud near \textbf{0.8~bars}, an \amhs~cloud formed from a reaction of \ammonia~and \hs~at \textbf{2.3 bars}, and a massive water ice/liquid solution cloud near \textbf{6~bars}, depending on assumptions of composition and thermal structure (see \citet{weidenschilling73}, \citet{atreya85}, and \citet{wong15}). Thermal emission from the deep atmosphere is attenuated by the variable opacity at 5~\um~ of one or more of these three cloud layers. The Great Red Spot has very low flux at 5~\um~ compared with adjacent regions. Previous 5-\um~ datasets such as Voyager IRIS, Galileo NIMS, and spectra from the Kuiper Airborne Observatory of Jupiter provided a wealth of information on Hot Spots (e.g. \citet{bjoraker86b, bjoraker86a} and \citet{roos04}) but they did not have the combination of sensitivity, high spectral resolution, and spatial resolution required to model the Great Red Spot. In this study, the use of high resolution instrumentation on the Infrared Telescope Facility (IRTF) and Keck telescopes in Hawaii allows us for the first time to characterize the cloud structure and gas composition of the Great Red Spot at pressures between 0.5 and 5~bars. The abundance of \water~in Jupiter's atmosphere is of fundamental importance in understanding the origin of Jupiter, the composition of its clouds, and jovian dynamics at pressures greater than 2~bars. One of the key objectives of the Juno mission, which began orbiting Jupiter in July 2016, is to measure water vapor below Jupiter's clouds to determine the O/H ratio using the Microwave Radiometer (MWR) \citep{janssen05}. Interpretation of these data may not be straightforward, however, due to the small microwave absorptivity of \water~ gas compared with \ammonia~ (see \citet{depater05} for details). Ground-based measurements of \water~ and \ammonia~ are important in order to provide upper boundary conditions to these key absorbers. Well-constrained values for these gases between 0.5 and 5~bars should improve the accuracy of the MWR's measurements of \water~ and \ammonia~ down to 100 bars, which only Juno can perform. In Fig.~1 we illustrate the complementarity of microwave and 5-\um~ spectra of Jupiter to probe the deep atmosphere. On the left, we show contribution functions \textbf{between 4.66~\um \ and 5.32~\um} \ due to gas opacity alone. The strongest absorption lines sound near 1 bar, while weak features can probe down to \textbf{6 or 7~bars}. Three clouds are shown at their predicted levels from thermochemical models (e.g. \citet{wong15}). On the right we illustrate weighting functions for each of the six channels of the Microwave Radiometer, adapted from \citet{janssen05}. There is excellent overlap in sounding Jupiter between 5~\um~ and 3 of the 6 MWR channels, namely at 3.125~cm, 6.25~cm, and 12.5~cm. Early Juno results on measurements of \ammonia~ in the deep atmosphere using MWR were presented by \citet{li17} and \citet{ingersoll17}. In July, 2017 Juno's orbit passed over the Great Red Spot yielding both spectacular images and microwave observations, which are currently being analyzed \citep{li17b}. \begin{figure*}[!ht] \begin{center} \begin{tabular}{ll} \hspace{-0.12in} \includegraphics[width=\textwidth]{NEBHot_TPcon4fig1.eps} \end{tabular} \vspace{-0.1in} \caption {\footnotesize { Contribution functions of Jupiter \textbf{between 4.66~\um \ and 5.32~\um} \ compared with MWR channels on Juno \citep{janssen17}. On the left, the Galileo probe temperature/pressure profile \citep{seiff98} is shown in red for comparison.}}\vspace{-0.1in} \label{fig1} \end{center} \end{figure*} In the next section we describe the instrumentation, observing circumstances, and data selection for the Great Red Spot. In Section 3 we present spectra of Jupiter's Great Red Spot from ground-based and Cassini data. We demonstrate that we can constrain the lower cloud structure of the GRS even when higher-altitude clouds greatly attenuate the thermal flux from the deep atmosphere. \citet{bjoraker15} modeled a Hot Spot in the South Equatorial Belt and a cloudy region in the South Tropical Zone. In the current study we apply the same methodology to constrain the pressure of the deepest clouds in Jupiter's Great Red Spot. In addition, we model ground-based spectra at 5.3~\um~ and Cassini spectra at 7.2~\um~ to constrain the location of upper and mid-level clouds in the GRS. Using a 3-layer highly-simplified cloud model we derive the abundance of \water, \ammonia, \ph, \germane, and CO in the deep atmosphere of the Great Red Spot. These abundances will help to understand the dynamics of the atmosphere in this unique feature on Jupiter. | \subsection{Cloud Structure} Using a simplified 3-cloud model, we obtained a good fit to the 5-\um~ spectrum of the Great Red Spot using a framework that is consistent with thermochemical models. Future studies can build on and test this framework by adding cloud microphysics, which would include modeling the vertical extent of each cloud layer. In addition, improved models would include wavelength-dependent absorption coefficients of \ammonia~ ice and \amhs~ ice across the 5-\um~ window, and incorporate scattering. \textbf{Our discovery of an opaque cloud at 5$\pm1$~bars (almost certainly a water cloud)} in the Great Red Spot may suggest that the vortex extends much deeper than the water cloud layer. \citet{marcus13} developed a model of secondary circulation within jovian anticyclones that maintains their key thermodynamic features: a high-pressure core with its associated geostrophic winds, sandwiched between a cold, high-density lid at the top of the vortex, and a warm, low-density anomaly at the base of the vortex. The cold lid in this scenario is a well-known observational feature \citep{cheng08, fletcher10}, and \citet{wong11} found that static stability measurements and models agreed with vortex models at the upper tropospheric levels where the cold lid is seen. \citet{wong11} also found that the low-density anomaly at the base of the vortex is consistent with stratification produced by condensation of a supersolar water cloud. If the GRS indeed had its base in the water cloud layer, then the warm, low-density anomaly there should generally inhibit water cloud formation, a scenario that is contradicted by our results in Figures 7 and 8. In fact, if the \citet{marcus13} model holds, our observation of 5-bar water clouds in the GRS may suggest that the vortex midplane lies below the 5-bar level, so that the water cloud layer is within the upper half of the vortex, where the secondary circulation is dominated by upwelling. Such a large vertical extent would also be consistent with preliminary Juno Microwave Radiometer observations of the GRS \citep{li17b}, as well as dynamical simulations showing that the vortex can extend well into the free convective zone \citep{chan13}. \textbf{Our derived pressure of the reflecting layer at 570$\pm30$~mbars in the GRS} is consistent with previous studies of upper clouds by \citet{banfield98}, but it is inconsistent with a layer at 900~mbars as suggested by \citet{simon01}. Our evidence for a thick middle cloud is consistent with near-infrared spectra from Galileo NIMS. \citet{irwin01} investigated the question of whether the \ammonia~ cloud or a deeper cloud is responsible for the observed variability in the brightness of Jupiter at 5~\um. They studied Galileo NIMS data and found an anti-correlation between the brightness of Jupiter at 5~\um~ and that at 1.58~\um. Irwin et al. used a reflecting layer model to fit NIMS spectra from 0.7 to 2.5~\um~ of belts (bright at 5~\um, dark at 1.58~\um) and zones (dark at 5~\um, bright at 1.58~\um). Successful models required variations in cloud opacity to lie deeper than 1 bar. Their model of the Great Red Spot from NIMS data required increased \ammonia~ cloud opacity to fit 0.7 to 1.0~\um~ as well as increased cloud opacity between 1 and 2 bars to fit the spectrum from 1.0 to 2.5~\um. Thus, Irwin et al. provided evidence for two cloud layers in the GRS, but they could not determine the optical depths of each cloud separately. The cloud structure derived for the GRS in this study is fairly similar to that derived for a cloudy feature in the South Tropical Zone (STZ) by \citet{bjoraker15}. In that paper we argued that mass flux into the upper cloud layers of the STZ is not dominated by horizontal transport, as suggested by \citet{showman05}, but instead is driven by vertical transport from below. Jupiter's circulation in zones (and now in the Great Red Spot) therefore maintains the same sign of upwelling/downwelling across the full 0.5-5~bar weather layer. The same sign of upwelling/downwelling over this large an extent was also derived by \citet{depater10} from 5-micron bright rings around vortices. These authors suggested that vortices must extend vertically from at least the 4 to 7-bar level up to the tropopause. This seems to be the case for the Great Red Spot as well. \subsection{Water abundance and the O/H ratio} This study constitutes the first detection of gaseous water in Jupiter's Great Red Spot. Water follows a saturated profile in the GRS, similar to that of a cloudy feature in the STZ, but unlike the highly depleted \water~ profile observed in Hot Spots \citep{bjoraker15}. The pressure that we derived for the opaque lower cloud is useful for understanding the vertical extent and dynamics of the GRS. In addition, it provides a constraint on the deep water abundance of Jupiter as a whole, and therefore its O/H ratio, assuming that the lower cloud is, in fact, composed of water. The base of the water cloud is sensitive to the deep abundance of water because higher abundances lead to condensation at deeper levels. Our data provide the level of the cloud top, not the cloud base. Since the top is at a higher altitude than the base, cloud top constraints provide lower limits to the pressure of the cloud base, or, lower limits to the deep abundance of water. The GRS spectrum requires a water cloud top at $P\ge5$~bar. The cloud base is therefore found at $P > 5$~bars. Following Fig.~1 in \citet{wong08}, a cloud base at 5 bar corresponds to an O/H ratio of 1.1$\times$ solar (corrected to the new solar O/H ratio of Asplund et al. 2009). The O/H lower limit of 1.1$\times$ solar is consistent with, and more constraining than, the Galileo Probe lower limit of 0.48$\pm$0.16$\times$ solar \citep{wong04}. Our observations support either solar or super-solar enrichments of water in Jupiter, providing somewhat better constraints to planetary formation models (\citet{gautier01}, \citet{hersant04}, and \citet{wong08}). \subsection{Ammonia abundance} \textbf{Our retrieved \ammonia~ abundance of 200$\pm50$~ppm in the Great Red Spot} for pressures greater than 0.7~bars may be compared with other 5-\um~ studies. \citet{giles17b} obtained 5-\um~ spectra of Jupiter using the CRIRES instrument on the Very Large Telescope at the European Southern Observatory in Chile. They did not observe the Great Red Spot, but they did report the latitudinal variation of \ammonia~ acquired using spectra aligned north-south on Jupiter's central meridian. They found NH$_3$ mixing ratios less than 200 ppm at 1.6 bars for all latitudes using a strong NH$_3$ absorption feature (at 1939~cm$^{-1}$; see Fig. 15), and larger values for NH$_3$ at 3.3 bars using a weaker feature (at 1929 cm$^{-1}$). Some differences between this group's work and our results may be due to the very different treatments of cloud opacity in the two studies. \citet{giles17b} initially used a single compact cloud layer at 0.8~bars. Later, they examined the effect of an additional cloud at 5~bars, following \citet{bjoraker15}. They also did not include sunlight reflected from an upper cloud layer. \citet{giles17b} obtained 150~ppm \ammonia~ at 1.6~bars increasing to 250~ppm at 3.3~bars for a region of Jupiter at 20\deg~ South latitude on Jupiter, which is the latitude of the GRS. This is in good agreement with our results. However, this agreement may be fortuitous if \ammonia~ varies with longitude due to local dynamics. Using microwave observations of Jupiter from the Very Large Array radio telescope, \citet{depater16} obtained longitude-resolved maps with sufficient spatial resolution (1200~km) to resolve the Great Red Spot. Using microwave channels from 5.49 to 17.38 GHz (1.7 to 5.5~cm), they obtained a vertical profile for \ammonia~ at the center of the GRS of 200~ppm at pressures less than 1.5~bars increasing to 570~ppm at deeper levels. Contribution functions for these wavelengths cover roughly 0.6 to 6~bars in Jupiter's troposphere, depending on the exact \ammonia~ profile (see Fig.~1). In a future paper we plan to model the GRS using an updated (and recalibrated) microwave spectrum between 5 and 25 GHz. We plan to include microwave opacity by clouds in addition to absorption by \ammonia, \hs, and \water~ (e.g., \citet{depater93, depater05}) for comparison with the current study at 5~\um. Our deep (0.7 - 5 bars) \ammonia~ concentration of 200~ppm is a low value, and has implications for the vertical profile of \ammonia~ in Jupiter's atmosphere both inside and outside of the GRS. Assuming that the internal circulation of the GRS is somewhat isolated from its exterior, the 200~ppm concentration of \ammonia~ is likely to persist all throughout the interior of the vortex, even below the water cloud that limits our observational sensitivity to pressures less than 5~bars. Our deep GRS value is lower than the concentration of about 400~ppm in the Galileo Probe entry site near 5 bars \citep{hanley09}, lower than the concentration of 570~ppm at the deepest levels of the Galileo Probe entry site \citep{wong04}, and lower than the deepest concentration of 360~ppm derived from Juno MWR data \citep{li17}. Measurements of the \ammonia~ profile in the Galileo Probe Entry Site, a 5-\um~ Hot Spot, found a gradually increasing concentration of \ammonia, down to a well-mixed level near 8~bars, coincident with a layer of higher static stability \citep{magalhaes02}. The 8-bar level possibly represented the ammonia cloud base \citep{wong09} after having been deflected downward by the Rossby wave system responsible for creating these Hot Spots \citep{showman98, friedson05}. This scenario implied that Jupiter's \ammonia~ concentration should be around 570~ppm everywhere below the cloud base, a high value that was difficult to reconcile with disk-averaged microwave spectra \citep{depater01}. But spatially resolved microwave spectroscopy can explain these discrepancies by showing that high (570~ppm) ammonia below the cloud base indeed occurs in plumes located near (and to the south of) 5-\um~ Hot Spots, while other regions of the planet are less ammonia-rich \citep{depater16}. This is qualitatively consistent with the Juno MWR findings of depleted (200~ppm) ammonia in the 3 - 10~bar range, over a wide range of latitudes [10 - 40$^{\circ}$ S and 20 - 40$^{\circ}$ N, (see \citet{li17}). Although \ammonia~ within the GRS should be vertically mixed by the internal circulation of the vortex, the composition of the air within the GRS should generally reflect the composition of its surroundings. \subsection{Phosphine abundance} Phosphine is a disequilibrium species in Jupiter's atmosphere. Equilibrium models by \citet{fegley94} show that \ph~ should be abundant at 1500~K in the deep atmosphere, but it would be converted to P$_{4}$O$_{6}$ at colder levels. The presence of \ph~ at the 5-bar level requires vertical transport from great depth on a time scale faster than the chemical conversion to P$_{4}$O$_{6}$ with an eddy diffusion coefficient (K) on the order of 10$^{8}$ cm$^{2}$ sec$^{-1}$. The value of K was constrained by measurements of \ph~ primarily in Hot Spots (e.g. \citet{bjoraker86b}). Our new measurements of \ph \ in the GRS indicate that vertical transport rates in the GRS may be higher than those in Hot Spots. \textbf{An alternate disequilibrium model was developed by \citet{wang16}. In this model \ph \ is converted to H$_{3}$PO$_{4}$ instead of P$_{4}$O$_{6}$. They found that \ph \ is relatively insensitive to K, and thus should not vary with location on Jupiter.} \textbf{Our \ph~ abundance in the GRS at the 5-bar level (0.8$\pm0.2$~ppm) is a factor of 2 higher than in the warm collar surrounding the GRS (0.4$\pm0.1$~ppm)}. \citet{bjoraker15} measured 0.45~ppm \ph~ in an SEB Hot Spot at 17\deg S and 0.7~ppm \ph~ in a region in the STZ at 32\deg S. These four measurements provide evidence for an enhancement in \ph \ at the 5-bar level in cloudy regions (GRS and STZ) compared with regions lacking an opaque water cloud (an SEB Hot Spot and the Warm Collar). This may be a consequence of enhanced convection in cloudy regions on Jupiter. \textbf{These measurements are not consistent with the predictions of the model developed by \citet{wang16}.} \citet{fletcher16} reported 1.2~ppm \ph~ over the Great Red Spot at 0.5~bars using IRTF/TEXES data at 10~\um. The \ph~ abundance fell off to 0.4 to 0.6~ppm in the warm collar surrounding the GRS. This falloff from the center of the GRS to the warm collar is in qualitative agreement with our measurements. However, our 0.8-ppm value for \ph~ at 5~bars is lower than the value at 0.5~bars (1.2~ppm) obtained using absorption features at 10~\um. \citet{fletcher09} first noticed the discrepancy between \ph~retrievals at 5 and 10~\um. A larger abundance of \ph~ at 0.5 bars than at 5~bars would be inconsistent with the disequilibrium model. These authors attributed this discrepancy to possible errors in the line strengths of \ph. \textbf{This discrepancy motivated a laboratory study of \ph \ line strengths at 5~\um \ \citep{malathy14}. The newer measurements of \ph \ line intensities are about 7\% \ higher than the older lab data of \citet{tarrago92} that are on the GEISA line atlas \citep{husson05} that we used in this study. This is not sufficient to explain the 50\% difference between \ph \ mole fractions derived from 5 and 10~\um \ data.} Thus, this problem has yet to be resolved. \subsection{Germane abundance} Germane is also a disequilibrium species in Jupiter's atmosphere. Models by \citet{fegley94} \textbf{and by \citet{wang16}} show that \germane~ should be abundant at 2000~K in the deep atmosphere, but it would be converted to GeS at colder levels. The presence of \germane~ at the 5-bar level requires vertical transport from great depth with an eddy diffusion coefficient similar to that obtained for \ph, $\approx$ 10$^{8}$ cm$^{2}$ sec$^{-1}$. \textbf{Our retrieved value of 0.35$\pm0.05$~ppb in the GRS} is compatible with this vertical transport model. \citet{giles17a} observed Jupiter's South Equatorial Belt (SEB) at 5~\um~ using the CRIRES instrument, as described above. They detected the strong Q-branch of the \nuthree~ band of \germane~ as well as the R3, R6, and R7 features. The R6 feature at 4.65~\um~ (2150.5~\wn) is the same feature that we observed in the Great Red Spot (see Fig.~20). Since the SEB does not have an opaque water cloud (see \citet{bjoraker15}) the pressure in the line formation region at 4.65~\um~ is higher than in the GRS. Therefore, absorption features due to various isotopologues of \germane~ are blended together in the SEB due to opacity broadening and pressure broadening by \h. By measuring \germane~ in a cloudy region such as the GRS, we can spectrally separate each individual isotopologue. In Fig.~19 and Fig.~20 we show absorption features of all 5 isotopic variants in the GRS. These are \gezero, \getwo, \gethree, \gefour, and \gesix. The terrestrial relative abundances of germanium isotopes are 20.6\%, 27.5\%, 7.8\%, 36.5\%, and 7.7\% respectively \citep{berglund11}. We obtained line lists for all isotopologues of \germane~ from R. Giles (private communication, see \citet{giles17a} for details). Using the same CRIRES dataset, \citet{giles17a} measured the latitudinal variation of the strong Q-branch of \germane~ on Jupiter. As with \ammonia, they found that it was difficult to separate the effects of spatially varying \germane~ from variations in the deep cloud structure. They found that they could fix \germane~ to 0.58~ppb, the value in the SEB, and allow the deep cloud to vary. Alternatively, they could allow \germane~ to vary with latitude, or they could permit both to vary. Their best fit model required both gas abundances and the deep cloud to vary with \germane~mole fractions ranging from 0.25 to 0.7~ppb. They obtained 0.45~ppb at 20\deg~ South, in good agreement with our results for the GRS. \subsection{Carbon monoxide abundance} There are two sources for CO on Jupiter. An external source of oxygen from meteoroids supplemented by large impacts such as from comet Shoemaker-Levy 9 results in the production of CO in Jupiter's stratosphere \citep{bezard02}. The second source of CO comes from \methane~ and \water~ in the deep atmosphere. Methane is the principal reservoir for carbon in Jupiter's reducing atmosphere, but at temperatures greater than 1000~K there is expected to be at least 1 ppb CO in chemical equilibrium with other carbon species including CH$_4$. Methane is the only equilibrium carbon species at the colder temperatures (273 K) probed at 5 $\mu$m. The unexpected detection of CO on Jupiter at 5~\um~ by \citet{beer75} led to the development of a disequilibrium model in which CO is transported from depth faster than the time scale to convert it back to \methane~ \citep{prinn77}. The Prinn and Barshay model was updated by \citet{bezard02} using improved reaction rates. The abundance of CO at 5~bars is proportional to the deep O/H ratio and depends on the eddy diffusion coefficient (K). Thus, the tropospheric abundance of CO can provide constraints on the deep water abundance as well as vertical transport rates. \citet{bezard02} measured 1.0 $\pm$ 0.2 ppb CO at 6~bars in an NEB Hot Spot. \textbf{Our abundance of 0.8$\pm0.2$~ppb CO in the GRS} is the first measurement of CO in a region of Jupiter with thick clouds. Our CO results, when combined with our measurements of \ph~ and \germane, suggest that vertical transport rates in the GRS are similar to or somewhat higher than those in Hot Spots. A value of 10$^{8}$ cm$^{2}$ sec$^{-1}$ for the eddy diffusion coefficient and 0.8~ppb CO would correspond to an O/H ratio of 4 times solar in the Bezard et al. (2002) model using the more recent solar abundances of \citet{asplund09}. \citet{bezard02} estimated that K lies in a range between 4x10$^{7}$ and 4x10$^{9}$ cm$^{2}$ sec$^{-1}$. With these large uncertainties in K, \citet{bezard02} were only able to constrain Jupiter's O/H ratio to be between 0.35 and 15 times solar (after correction to the solar value of \citet{asplund09}). Using our lower limit to O/H of 1.1 times solar derived from the pressure of the water cloud, we can narrow the range of the O/H ratio to be between 1.1 and 12 times solar. \textbf{\citet{wang15} proposed a new formulation for Jupiter's eddy diffusion coefficient. They investigated the range of water enrichment required to fit Jupiter's CO abundance using two different CO kinetic models. Using kinetics from \citet{visscher11}, and a CO abundance of 1~ppb from \citet{bezard02}, they inferred O/H enrichments ranging between 0.1 and 0.75 times solar. Using an alternate chemical kinetics model originally applied to hot Jupiters \citep{venot12}, they constrained the O/H ratio to be between 3 and 11 times solar. Our lower limit to O/H of 1.1 times solar derived from the pressure of the water cloud is consistent with the \citet{venot12} kinetic model, but not with the model of \citet{visscher11}. Using our CO value of 0.8~ppb in the Great Red Spot and the \citet{venot12} kinetic model, we obtain O/H enrichments between 2.4 and 9 times solar}. | 18 | 8 | 1808.01402 |
1808 | 1808.10602_arXiv.txt | The Effective Field Theory (EFT) of Preheating with scalar fields, implies three types of derivative couplings between the inflaton and the reheating field. Two of these couplings lead to scales below which only one of the two species appear as the low energy modes. In this paper, the variety of low energy regimes in terms of the species they accommodate are explored by studying the scales introduced by the derivative couplings and the dispersion relations they lead to. It is noted that the EFT of two scalar fields can give rise to non-trivial sound speed for both the inflation and reheating sector even at scales where modes of both species propagate freely, suggesting the presence of additional heavy fields. The regimes where one of the species affects the dispersion relation of the other while not appearing as an effective mode itself, are named as ``Hidden Regimes'' during preheating. | In single field inflation, only one type of field, the inflaton, dominates the overall energy momentum density of the universe. At the background level, this field has some time dependence which leads to a time dependent cosmological background $H(t)$. This background will not remain invariant under time diffeomorphisms. However the time diffeomorphisms are a symmetry of the action. On such a time dependent background, there exists modes of a physical scalar perturbation which transforms nonlinearly under further time diffeomorphisms, making sure that the action itself remains invariant. These scalar modes are the inflationary scalar perturbations, observed as temperature fluctuations in the Cosmic Microwave Background (CMB). They lead to formation of protogalaxies in the early universe, on an otherwise homogeneous background. The present day universe is filled with many other types of matter, such as the particles of standard model, dark matter and dark energy. These correspond to perturbations of different types of matter fields, not all of which are scalar. It is an open question to understand how and when these fields start playing an important role in the universe overtime. It could be that a number of fields, nontrivial kinetic terms or non-minimal couplings are responsible for the underlying mechanism of inflation itself. While these cases are being more and more constrained by observations within a window of time during inflation \cite{Akrami:2018odb}, trying to understand couplings of the inflaton to other species is what is at hand for understanding the passage to other matter sources that are present in today's universe. It is likely that perturbations of noninflationary species begin to populate the universe after the end of Inflation rather then during \cite{Armendariz-Picon:2017llj}. The period of inflation corresponds to an accelerated expansion, with a weakly time dependent background $H_I(t)$ that resembles the approximate de Sitter spacetime. Single field inflation ends when the Hubble parameter drops below the mass of the inflaton field $m_\phi$, $H_I(t_{end})\leq m_\phi$. At times of order $m_\phi t\gg 1$, the Firedmann equations give a matter dominated solution with oscillatory corrections whose frequency is set by the mass $m_\phi$. This era is referred to as preheating, during which the time dependence of the background $H_p(t)$ is strong. This time dependence works into the couplings of the inflaton field with perturbations of other species, and leads to their resonant growth leading to a more efficient production of the reheating modes compared to perturbative decay alone \cite{Kofman:1997yn}. At the end of this intermediate stage, the inflaton will decay into fields lighter then itself, giving rise to a radiation dominated phase. In \cite{Ozsoy:2015rna, Giblin:2017qjp}, general interactions for the perturbations of the inflaton and a scalar reheating field were studied during preheating with effective field theory (EFT) methods. It was noticed that studying these perturbations during preheating without addressing the dynamics that can give rise to the background $H_p(t)$ in the fashion of the EFT of quasi single field inflation \cite{Noumi:2012vr}, gives more insight then an EFT for the dynamics of the two species, which have been rather useful in the case of inflation \cite{Weinberg:2008hq,Assassi:2013gxa}. In this EFT approach, being associated with the species that drives the time dependence of the background, scalar inflationary perturbations $\delta g^{00}$ are Goldstone modes that nonlinearly realize time diffeomorphism invariance while reheating sector is introduced as an unspecified scalar perturbation $\chi$. This difference in their nature, leads to different interactions for the different species. How the background behavior $H_p(t)$ enters the quadratic self couplings of each species determines the resonant particle production, and the hierarchy $H_p(t)\ll m_\phi$ between the background scales, leads to hierarchies among the scales important for the dynamics of the perturbations \cite{Giblin:2017qjp}. In addition to interactions that lead to particle production, the EFT of preheating involves three different types of derivative couplings among the inflationary and reheating perturbations. The focus of this work are these derivative interactions. In the first work to address derivative couplings of the inflaton and a scalar reheating field, through an analysis of instability bands it was observed that the derivative coupling of interest does not promise very efficient production of reheating modes \cite{ArmendarizPicon:2007iv}. Among the three derivative EFT couplings, two of them describe energy scales. On a complimentary line to the conventional analyses of instability bands, the main effort of the present work is to understand which of the two species of perturbations occur as the effective degree of freedom at what scales, by considering non relativistic limits and using the methods of \cite{Baumann:2011su}. The conclusion is that, with interactions of the type considered in section \ref{sec:beta1}, the effective lowest energy modes are the inflaton modes. For these type of couplings, the reheating perturbations work to adjust the dispersion relation of the inflaton rather then being present themselves. This suggests that the role of the reheating sector here is to assist the inflationary sector rather then being likely to be produced. The low energy regime of these couplings are named as the regime of Hidden Preheating, in the sense that the presence of the reheating field is hidden. The original example of derivative couplings falls into this category. The situation is reversed with the reheating perturbations being the low energy modes whose dispersion relation is adjusted by the inflaton modes in the presence of couplings considered in section \ref{sec:beta3}. However the couplings of sections \ref{sec:beta2} and \ref{sec:beta3} give rise to further modifications in the dispersion relation that exist even at scales where the inflaton and reheating modes propagate freely. Similar to the polynomial couplings of light fields to heavy ones giving rise to corrections to the mass of the light fields at low energies, derivative couplings give rise to corrections for the dispersion relation of light modes at low scales. This suggests that these later kind of couplings signal the presence of more fields that are actually present and interact with the inflaton and the reheating sectors, but who are themselves too heavy to appear as propagating degrees of freedom.% The text is organized as follows. The Effective Field Theory formalism of cosmological perturbations and how it captures preheating is reviewed in section \ref{sec:EFTreview}. Section \ref{sec:background scales} together with appendix \ref{ap:background} review the general properties of the preheating background $H_p(t)$ and scales associated with important processes in the EFT. The main consequences of the derivative couplings are explored in section \ref{sec:Hidden} and section \ref{sec:conc} summarizes the main results. | \label{sec:conc} Low energy effective field theories (EFT), especially the ones that are developed at the level of perturbations such as the EFT considered here, aim to capture the variety of interactions in the most general way. This generality is achieved by considering all of the interactions allowed by the symmetries that are present below a specified scale. This scale in the EFT set up considered here was the scale of spontaneous breaking of time translation invariance, due to the time dependent nature of the background $H(t)$. Among the possible interactions for the inflaton and reheating field perturbations, the present work has focused on the extra derivative couplings that appear under three different classes, specified by the EFT parameters $\{\beta_1,\beta_2,\beta_3\}$. The scales these derivative couplings introduce, the nature of the effective degrees of freedom at energies below the introduced scales ( whether they are inflaton perturbations or the scalar reheating perturbations), and the corrections to the dispersion relation for the effective modes at low energies have been explored in this work. The properties of the background as determined by the preheating era, that is the presence of two scales $m_\phi$ and $H_m$ with the hierarchy $m_\phi\gg H_m\simeq H_p$ between them, have been used to determine the hierarchy between the scales of the interactions, such as the particle production scale and the scale associated to derivative couplings. All this has led to the main conclusion that, all though the aim of preheating is to capture energy transfer between two different species, here the inflaton and the scalar reheating sector $\chi$, in the presence of such derivative couplings only one of the species propagates as an effective degree of freedom at very low energies, while the other stays hidden and modifies the dispersion relation of the propagating species. Instead of an analysis of instability bands to determine the efficiency of $\chi$-production, the main pursuit here has been the identification of the relevant species for low energies and exploring how the dispersion relation of this species gets modified. It is left for future to discuss the efficiency in production of the identified low energy modes through studying the details of resonance in comparison to perturbative decay rates. While this EFT method allows one to study the properties of perturbations right away, the disadvantage can be that it is not always clear what kind of interactions at the background level would give rise to these interactions at the level of perturbations. For example, an interaction of type $(\frac{1}{\Lambda}\partial^\mu\phi\partial_\mu\phi)X$ where $\phi$ is the inflaton which is to be expanded as $\phi(\vec{x},t)=\phi_0(t)+\delta\phi(\vec{x},t)$, $X$, with $X(\vec{x},t)=X_0(t)+\chi(\vec{x},t)$, is the reheating field and $\Lambda$ is some mass scale; is an example that gives rise to $\beta_1$ type couplings. And indeed these type of interactions are common in inflationary literature in many studies that wish to respect the shift symmetry for the inflaton. On the other hand, $\frac{1}{\Lambda}(\partial^\mu X\partial_\mu X)\phi$ would be an example to $\beta_3$ type couplings, which however would not be an interaction to consider if one is concerned with a shift symmetric inflaton. Looking at the preheating literature, the interactions considered are more of a polynomial type, for instance $g^2\phi^2X^2$ is the first case that has been considered. Derivative couplings during preheating have not been studied at the level they have been during inflation. So far, derivative couplings in preheating literature involve examples of only the class of $\beta_1$ couplings, among the three classes that the EFT methods suggest. Moreover the examples to $\beta_1$ case that have been studied are noted to be not very efficient for resonant production of low energy reheating modes, $\tilde{\chi}_{ck}$. Looking at the dispersion relations, here it is noted that at scales below the scale of derivative coupling $R_1$, the reheating modes appear to effect the canonical momenta of the inflaton perturbations, leaving them as the low energy species with a modified dispersion relation derived in equations \eqref{hiddenbeta1_rkm} and \eqref{omega-smallest}. On the other hand, $\beta_3$ interactions accommodate the reheating modes as the light degrees of freedom with a modified dispersion relation of \eqref{hiddenregime_beta3_dis}. Hence these later type of interactions may be more promising for resonant production of $\chi$ through derivative couplings. Moreover some of the derivative couplings, in the presence of $\beta_2$ and $\beta_3$ imply a sound speed and modified dispersion relations for both of the species even at energies where modes of both propagate freely. This suggests that these EFT coefficients may address models that involve additional heavy degrees of freedom. The reheating sector $\chi$ as considered here is quite general. Being a primordial scalar field, $\chi$ is most likely to contribute to structure formation and resemble fields associated to dark matter. In principal any of these couplings can arise in models of multi-field inflation. Since the effective field theory method at the level of perturbations followed here considers all possible terms that respect the symmetries at the scales of interest, one of the expected benefits of this approach is to come across new types of interactions that may not have been thought of yet. The appearance of the less explored case of $\beta_3(t)$ type couplings are an example to this. They would arise from attempts of generalizing interactions of the reheating field, as sketched in the previous paragraphs. It is left for future to explore for this later case, the phenomenological implications and the detailed structure of resonance in comparison with rate of perturbative decays. With regards to perturbative decay rates, the EFT interactions would give the possible Feynmann diagrams to be computed, however the strength of the amplitude from these diagrams will depend on the coupling parameter which in turn depends on the background physics. The background information is determined by how $H_p(t)$ and $m_\phi$ work into $\phi_0(t)$ and $X_0(t)$. The same holds for the efficiency of resonant particle production. One can make estimates on the scale of particle production from the general behavior of the background as it has been done here, but to study the actual efficiency one again has to first study the details of how the background parameters work into the Mathieu variables. Once solid examples that give rise to $\beta_3$ type interactions at the level of perturbations are constructed, then how the coupling parameters depend on $m_{pl},$ $H$ and $m_\phi$ through the background behavior of $\phi_0$ and $X_0$ will become more clear and, the perturbative decay rates and efficiency of resonance can be studied more concretely. % | 18 | 8 | 1808.10602 |
1808 | 1808.02221_arXiv.txt | For seven decades, the widely held view has been that the formation, the migration and the decay of short-lived starspots explain the constantly changing light curves of chromospherically active stars. Our hypothesis is that these deceptive observed light curves are interference of two real constant period light curves of long-lived starspots. The slow motion of these long-lived starspots with respect to each other causes the observed light curve changes. This hypothesis contradicts the current views of starspots. Therefore, we subject it to eight \hattuuslippa ~tests. Our new period finding method detects the two real light curves of FK Com. Our hypothesis is a total success: all real light curve parameters are directly connected to the long-lived starspots which are also seen in the Doppler images of FK Com. These parameters are spatially and temporally correlated just like in the Sun, including weak solar-like surface differential rotation. As for other chromospherically active stars, all eight \hattuuslippa ~tests also support our hypothesis. It explains many spurious phenomena: the rapid light curve changes, the short starspot life-times, the rapid rotation period changes, the active longitudes, the starspot migration, the period cycles, the amplitude cycles and the minimum epoch cycles. It also explains why the light curves and the Doppler images give contradicting surface differential rotation estimates even for the same individual star, as well as the abrupt 180 degrees shifts of activity (the flip-flop events) and the long-term mean light curves. We argue that the current views of starspots need to be revised. | \label{Intro} The ancient Egyptian papyrus Cairo 86637 is the oldest preserved document of the discovery of a variable star, Algol \citep[][]{Por08,Jet13,Jet15,Por18}. Algol's changes can be observed with naked eye, but the solar luminosity changes only with satellites \citep{Wil91,Rad18}. John of Worcester made the oldest preserved drawing of a sunspot in the year 1128 \citep[][]{van96}. \citet{Sch44} discovered the eleven years cycle in the number of sunspots. \citet{Hal08} discovered the Zeeman effect of solar magnetic field, and that this magnetic field is stronger in the sunspots. \citet{Kro47} discovered the starspots in the light curves of the eclipsing binary AR Lac. He observed short-term % changes ``within a few weeks to a few months''. FK Comae Berenices (HD 117555, FK Com) was among the first late--type stars where the starspots were also discovered \citep{Chu66}. This chromosperically active single G4 giant \citep{Str09} is the prototype of a class of variable stars, the FK~Com class, defined by \citet{Bop81} as rapidly rotating single G--K giants. Only a few stars belonging to this class have been found \citep[][]{Puz14,How16,Puz17}. These stars may represent recently coalesced W~UMa binaries \citep{Web76,Egg89,Wel94}. % The starspots of FK Com seem to concentrate on two long--lived active longitudes separated by 180 degrees, and undergo abrupt shifts between these longitudes. These shifts are called the \flip ~events \citep{Jet91,Jet93}. They occur also in binaries \citep[e.g.][$\sigma$ Gem]{Jet96}. Kuiper method analysis of the light curve minimum epochs gave the $P_{\mathrm{act}}=2.^{\mathrm{d}}401155 \pm 0.^{\mathrm{d}}000092$ active longitude period of FK Com \citep{Hac13}. There are numerous photometric \citep[e.g.][]{Kor02,Ola06,Pan07} and Doppler imaging \citep[e.g.][]{Kor00,Kor04,Kor07,Hac13,Vid15} studies of FK Com. % \citet{Jet17} presented a general light curve model for the Chromospherically Active Binary Stars\footnote{Table \ref{abbre} gives the abbreviations used in this paper.} (hereafter CABS). They studied the long-term Mean Light Curve (hereafter MLC) of fourteen CABSs. \citet[][JHL = Jet\-su, Henry, Leh\-ti\-nen]{Jet17} presented the hypothesis \begin{description} \item[-] \JHLhyp: \JHLtext \end{description} At first sight, this hypothesis would seem to contradict the current views of starspots and stellar Surface Differential Rotation (hereafter \sdr). Two methods are widely used to measure SDR \citep[][Sect. 7]{Str09}. In the \lcmethod, the $P_{\mathrm{phot}} \! \approx \! P_{\mathrm{rot}}$ estimates are obtained from period analysis \citep[e.g.][]{Rei13,Leh16} or spot modelling \citep[e.g.][]{Kip12,Lan14} of light curves. In the \dimethod, Doppler imaging gives $P_{\mathrm{rot}}$ periods of starspots at different latitudes \citep[e.g.][]{Pet04,Col07,Kov17}. These two methods can give {\it different} \sdr ~estimates for the {\it same} star, e.g. the \dimethod ~study by \citet{Kor00} indicated ``solid body rotation'' in FK Com, while \citet{Hac13} measured photometric period changes of 3.1\% with the \lcmethod. If the \JHLhyp ~is true, it should predict the results for the following eight undermining tests. \begin{description} \item[-] \Testonetext \item[-] \Testtwotext \item[-] \Testthreetext \item[-] \Testfourtext \item[-] \Testfivetext \item[-] \Testsixtext \item[-] \Testseventext \item[-] \Testeighttext \end{description} \noindent We summarize the results for all tests in Sect. \ref{Conclusions}. | \label{Conclusions} Our \JHLhyp ~is a total success, because we detect the two real light curves of FK Com with our new period analysis method. All parameters of these real light curves are directly connected to the long-lived starspots of FK Com. These starspot parameters are spatially and temporally correlated just like in the Sun, except that dark starspots change the luminosity of FK Com (Sect. \ref{FKsdr}: Eq. \ref{FKmode}). The simultaneous Doppler images confirm weak solar-like surface differential rotation. Since FK Com resembles the Sun and holds the rotation record for single stars \citep[][``King of spin'']{Ayr16}, its starspots provide valuable observational and theoretical constraints for the starspots of other slower rotating chromospherically active stars. We present an analogy of incompatibility, because for seven decades these ``creatures'' have managed to evade detection behind the ``veil'' of interference. Imagine a face with a left and a right eye. Both eyes can disappear and reappear. At any moment, the number of eyes may be zero, one or two. The original stationary right eye can disappear and reappear only at fixed locations. The original non-stationary left eye rotates slowly around the head. Both eyes also wobble up and down simultaneously (Fig. \ref{FKspots}: Alternative \ref{sdrone}). We see this head spinning. Soon it is impossible to tell which eye is the original left or right eye (map-\Inc ~after $P_{\mathrm{lap}}/2$). The only compatible pictures of this face are snapshots, but these snapshots can not be used to recognize this constantly changing face. These snapshots can capture only one side of the head, or equivalently half of the full visible surface of FK Com. As for other chromospherically active stars, this ``spinning head'' hypothesis \begin{description} \item[-] \JHLhyp: \JHLtext \end{description} \noindent is subjected to eight \hattuuslippa ~tests. We use the notations $P_1$ and $P_2$ for periods of the $g(t)$ model (Eq. \ref{fullmodel}). They fulfill $P_1<P_2$. Depending on the object, these $P_1$ and $P_2$ periods can refer to the $P_{\mathrm{act}}$, $P_{\mathrm{rot}}$ or $P_{\mathrm{orb}}$ periods. The {\it binaries} may have $P_1=P_{\mathrm{act}}$ and $P_2=P_{\mathrm{rot}} \approx P_{\mathrm{orb}}$, or $P_1=P_{\mathrm{rot}} \approx P_{\mathrm{orb}}$ and $P_2=P_{\mathrm{act}}$. The {\it single} stars may have $P_1=P_{\mathrm{act}}$ and $P_2=P_{\mathrm{rot}}$, or $P_1=P_{\mathrm{rot}} $ and $P_2=P_{\mathrm{act}}$. For the {\it binaries}, the relation $P_{\mathrm{rot}} \approx P_{\mathrm{orb}}$ can be used to check, if $P_1$ or $P_2$ represents the stationary $P_{\mathrm{rot}}$ period. This is not possible for the {\it single} stars, like FK Com. Here are our eight tests and our results for these tests: \begin{description} \item[-] \Testonetext \item[] \result{ One-dimensional period finding methods ``detect'' {\it many different spurious} periods for the {\it same} star, although the {\it real} $P_1$ and $P_2$ periods remain the {\it same} (Sect. \ref{periodinc}).} \item[-] \Testtwotext \item[] \result{The Kuiper method analysis of the incompatible epochs did not give an unambiguous $P_{\mathrm{act}}$ period (Sect. \ref{minimainc}).} \item[-] \Testthreetext \item[] \result{This method does not detect only the $P_{\mathrm{act}}$ period. The detected period can be the $[(f_1+f_2)/2]^{-1}$, $f_2^{-1}$ or $f_1^{-1}$ period of $g(t)$ model, or none of these, because the amplitudes of the real light curves change (Sect. \ref{minimainc})} \item[-] \Testfourtext \item[] \result{ \citet{Jet17} detected the long-term mean light curves following the $P_{\mathrm{rot}}\approx P_{\mathrm{orb}}$ period. Due to minima-\Inc, their $P_{\mathrm{act}}$ estimates failed (Sect. \ref{minimainc}).} \item[-] \Testfivetext \item \result{{\it Interference} causes these events. The {\it abrupt} shift is $\Delta \phi_{\mathrm{b}}= 0.5\equiv 180\degr$, if the real light curves are equal amplitude sinusoids with constant periods $P_1$ and $P_2$ (Eq. \ref{abruptphase}).} \item[-] \Testsixtext \item \result{\citet{Oza18} have detected the stationary $P_{\mathrm{rot}} \approx P_{\mathrm{orb}}$ period and the non-stationary $P_{\mathrm{act}}$ period in the spot models for light curves of KIC~11560447 (Fig. \ref{ozafig}).} \item[-] \Testseventext \item[] \result{{\it Particular} and {\it general} evidence indicate that these periods have been detected in the Doppler images (Sect. \ref{SIsect}).} \item[-] \Testeighttext \item[] \result{The one-dimensional period finding methods detect {\it spurious} periods from the light curves. The Doppler images can detect {\it real} periods.} \end{description} The \JHLhyp ~neatly explains many phenomena that have been detected earlier from the light curves with one-dimensional period finding methods \begin{description} \item[-] Spurious {\it observed} rapid light curve changes \item[-] Spurious short starspot life-times \item[-] Spurious rapid photometric rotation period changes \item[-] Spurious active longitudes \item[-] Spurious starspot migration \item[-] Spurious cycles in periods and amplitudes \item[-] Spurious \flip ~events and \flip ~cycles \item[-] Long-term mean light curves \item[-] \lcmethod ~and \dimethod ~discrepancies for \sdr \end{description} \noindent While the above list of phenomena is by no means complete (e.g. Sect. \ref{minimainc}: light curve {\it shape}), it does contain at least the most obvious ones. It took a quarter of a century to find the correct result for the \Testfive: ``Farewell \flip.'' | 18 | 8 | 1808.02221 |
1808 | 1808.02367_arXiv.txt | {Among the intermediate mass, magnetic chemically peculiar (MCP) stars, CU~Vir is one of the most intriguing objects. Its 100\,\% circularly polarized beams of radio emission sweep the Earth as the star rotates, thus making this strongly magnetic star the prototype of a class of non degenerate stellar radio pulsars. While CU~Vir is well studied in radio, its high-energy properties are not known. Yet, X-ray emission is expected from stellar magnetospheres and confined stellar winds.} {Using X-ray data we aim to test CU Vir for intrinsic X-ray emission and investigate mechanisms responsible for its generation.} { We present X-ray observations performed with {\it XMM-Newton} and {\it Chandra} and study obtained X-ray images, light curves and spectra. Basic X-ray properties are derived from spectral modelling and are compared with model predictions. In this context we investigate potential thermal and non-thermal X-ray emission scenarios. } {We detect an X-ray source at the position of CU~Vir. With $L_{\rm X} \approx 3 \times 10^{28}$~erg\,s$^{-1}$ it is moderately X-ray bright, but the spectrum is extremely hard compared to other Ap stars. Spectral modelling requires multi-component models with predominant hot plasma at temperatures of about $T_{\rm X}= 25$~MK or, alternatively, a nonthermal spectral component. Both types of model provide a virtually equivalent description of the X-ray spectra. The {\it Chandra} observations was performed six years later than the one by {\it XMM-Newton}, yet the source has similar X-ray flux and spectrum, suggesting a steady and persistent X-ray emission. This is further confirmed by the X-ray light curves that show only mild X-ray variability. } {CU~Vir is also at X-ray energies an exceptional star. To explain its full X-ray properties, a generating mechanism beyond standard explanations like the presence of a low-mass companion or magnetically confined wind-shocks is required. Magnetospheric activity might be present or, as proposed for fast rotating strongly magnetic Bp stars, the X-ray emission of CU~Vir is predominantly auroral in nature. } | The A0p star CU Vir (HD 124224, HR 5313) is an enigmatic star of the upper part of the main sequence located at a distance of about 79~pc \citep{lee07}. With its $V= 5.0$~mag it is among the prominent, nearby, magnetic, intermediate mass, chemically peculiar stars (MCP star, ApBp star) that shows pronounced photometric and spectroscopic variability \citep[see e.g.][]{bab58}. Furthermore, it is so far a unique main sequence star that shows regular radio pulses persisting over decades, resembling the radio lighthouse of pulsars and interpreted as auroral radio emission similar to those observed on planets \citep[see e.g.][]{tri00, tri11}. In contrast to the rotationally modulated gyrosynchrotron radio emission commonly observed in MCP stars, the 100\,\% right-handed circularly polarized radio pulses from CU~Vir are explained by the electron cyclotron maser emission (ECME) mechanism. The polarization sense indicates that these originate from the northern polar regions of the oblique stellar magnetosphere, where annular rings emit narrow radio beams that sweep over the Earth location twice per stellar rotation. With a period of about $P_{\rm rot} = 0.52$~d, CU Vir is an unusually fast rotator for its class that shows in addition alternating variability of its rotation period over decades \citep{mik11,krt17}. Beside numerous radio studies, several dedicated multiwavelength campaigns on CU Vir were executed. A detailed study of its variability in optical/UV emission and its spectral energy distribution showed that this variability can be explained by strong spots of elemental over- and under-abundances \citep{krt12}. A mapping of CU~Vir's abundance anomalies and magnetic field based on spectropolarimetric observations as well as a reassessment of the stellar parameters was recently performed in \cite{koch14}. According to this study, CU~Vir is viewed under an inclination of about $i = 46^{\circ}$; its magnetic dipole axis is tilted by about $\beta =79^{\circ}$ to the rotational axis and has a field strength of $B_{d}= 3.8$~kG. The derived magnetic maps show a dipolar-like field topology that is non-axisymmetric with large differences between regions of opposite polarity, providing a natural explanation for the north-south asymmetry in the observed radio pulses. The spectral classification of CU~Vir (B9p\,/\,A0p) as well as its stellar parameters vary a bit in literature, recent values as determined by \cite{koch14} are $M=3.1 M_{\odot}$, $R=2.1 R_{\odot}$ and $T_{\rm eff}=12750$~K. Using $Vsini=145$~km\,s$^{-1}$ this implies with $i = 46^{\circ}$ a rotational speed of about 200~km\,s$^{-1}$. While rotating fast, probably due to its youth, CU~Vir is still nearly spherical and effects of gravitation darkening are moderate with temperature contrast of only a few 100~K. The existence of intrinsic X-ray emission from late-B and early-A stars has been debated, as 'normal' late-B\,/\,early-A stars are expected to be virtually X-ray dark, since these stars neither drive magnetic activity nor strong stellar winds. Often, low-mass coronal companions are suspected and identified as the true X-ray source. However, intrinsic X-ray emission from magnetically confined wind shocks (MCWS) could be also expected in the Ap/Bp stars. In this model, originally proposed to explain the X-ray emission of the A0p star IQ~Aur \citep{bab97}, the stellar wind from both hemispheres is channelled by the magnetic field and collides in the vicinity of the equatorial plane, leading to strong shocks and thereby plasma heating to X-ray temperatures. Advanced MCWS models using magneto-hydrodynamic simulations (MHD) or rigid rotating magnetospheres models (RRM) are now 'standard' models and have been used to interpret the observed X-ray properties of magnetic massive stars \citep[e.g.][]{dou14}. A 3-D model able to simulate the the incoherent gyrosynchrotron radio emission from a typical rapidly rotating magnetosphere of a hot magnetic star has been previously developed \citep{tri04}. Using this model, the multiwavelength radio light curves of CU~Vir have been simulated, constraining the magnetospheric physical conditions \citep{leto06}. Recently, the same simulation approach has been successfully applied also to the cases of the fast rotating and strongly magnetic B2Vp stars HR7355 and HR5907 \citep{leto17,leto18}. These two stars show evidences of non-thermal X-ray emission. By using the 3-D model, computed to simulate the gyro-synchrotron radio emission of these two stars, it was shown, that in addition to thermal plasma heated by the shocked magnetically confined wind streams, a nonthermal auroral X-ray radiation is also expected. Assuming a common framework able to explain the radio emission features of the hot magnetic stars, it is expected that similar auroral X-ray emission shall be observable also from CU~Vir. The star CU Vir was observed by {\it XMM-Newton} in 2011 and detected in X-rays \citep{rob16}, but the combination of positional offset and fuzziness of PSF made a confirmation of the detection desirable. For this purpose a {\it Chandra} ACIS observation was initiated by us and performed in 2017. Here we present a detailed analysis of the available X-ray data from CU Vir and put it into context of current models to explain the X-ray emission in magnetic intermediate mass stars. | \label{dis} The derived X-ray properties of CU~Vir and potential X-ray generating mechanisms are discussed in the following. \subsection{Intrinsic vs. extrinsic emission} The sharp {\it Chandra} PSF and good positional match make a chance alignment with an unrelated object very unlikely. This is specially true at its galactic latitude of +58.6~deg and its X-ray spectra without any hints on absorption on the line of sight and requiring multiple thermal or thermal+nonthermal components. In contrast, a very close and so far undiscovered companion to CU~Vir is not ruled out by our data. Given the measured X-ray luminosity of $\log L_{\rm X} = 28.4$, a late-type star is a viable option. However, the extraordinarily hard X-ray spectrum argues against this hypothesis for our target. Even active M dwarfs emitting a several times $10^{28}$~erg\,s$^{-1}$ in X-rays, i.e. those that are comparable or moderately X-ray brighter than CU~Vir, have significantly cooler coronae with average temperatures of 6\,--\,8~MK and spectral energy distributions that peak around 8~MK \citep[e.g.][]{rob05}. Stars that possess coronae with average temperatures at quasi-quiescent level as observed for CU~Vir exist, however these have so far only been observed in stars with X-ray luminosities around or exceeding the $\log L_{\rm X} = 30$ level, i.e. objects that are hundred times X-ray brighter. These luminosities are ruled out, as the distance to CU~Vir and thereby flux to luminosity conversion, is well established. The recent Gaia DR2 gives $d= 72 \pm 2$~pc, even reducing the above luminosities by about 20\,\%. The shape of the observed X-ray light curves, each obtained over several tens of ks, also clearly argues against an origin in a very strong flare that could produce sufficient amounts of correspondingly hot plasma. Furthermore, the repeated X-ray detection at a similar brightness and spectral hardness in observation separated by several years suggests that the derived X-ray properties reflect the typical state for CU~Vir. Overall, coronal emission from a stellar companion might contribute, but is an unlikely explanation for the bulk of X-rays that are observed from CU~Vir. As the possibility of a chance alignment with an extragalactic or galactic counterpart is negligible, intrinsic mechanisms that are capable of generating the observed X-ray emission are considered to be the most likely explanation for the X-ray detection. \subsection{CU Vir in context of magnetic intermediate mass stars} The strong violation of both stellar scaling relations by CU~Vir highlights its outstanding character also at X-ray energies and in combination with its moderate X-ray luminosity and hard X-ray spectrum it is so far unique in the late-B/early-A star regime. Although several stars in the MCP sample presented in \cite{rob16} have similar characteristics in at least one or two of the relevant parameters, indicating similarities in their X-ray generating mechanism, none is as extreme as CU~Vir. Even typical magnetic early-B stars with similar $\log L_{\rm X}/L_{\rm bol}$ typically have significantly cooler/softer X-ray emission \citep{osk11}. Nevertheless, comparably hard X-ray spectra combined with high $\log L_{\rm X}/L_{\rm bol}$ are also known for a few magnetic early-B stars like HD~182180/HR~7355 and HD~142184/HR~5907 \citep{naze14,leto17, leto18} or $\sigma Ori$~E \citep{san04b}, however these have X-ray luminosities that are by factors of several tens or even hundreds above the one of CU~Vir. Intriguingly, also the X-ray spectral properties are quite similar between CU Vir and the fast rotating Bp stars HR~7355 and HR~5907 \citep{leto17, leto18}. Albeit, the X-ray luminosity of these Bp stars is one to two orders of magnitude higher, the basic X-ray spectral properties and X-ray activity as well as X-ray/radio ratios are comparable. A discussion relating to the auroral model proposed for these stars is given in the next section. In Table~\ref{comp} we compare relevant parameters to highlight their similarities and differences. \begin{table}[t] \begin{center} \caption{\label{comp}CU Vir vs. early Bp stars} \begin{tabular}{lccc}\hline\\[-3mm] & CU Vir & HR 5907 & HR 7355\\\hline\\[-3mm] $M_{*}$ [M$_{\odot}$] & 3.1 & 5.5 & 6.0 \\ $P_{\rm rot}$ [d] & 0.52 & 0.51 & 0.52\\ $B_{\rm p}$ [kG] & 3.8 & 15.7 & 11.6\\\\[-3mm] $\log L_{\rm X}$ [erg\,s$^{-1}$] & 28.4 &30.1 & 30.0\\ kT/$\alpha$ & 0.9/2.0 & 1.0/1.6 & 1.0/1.7 \\ $\log L_{\rm X}/L_{\rm bol}$ & -7.1 & -6.4 & -6.5\\ log $L_{\rm X}/L_{\rm rad}$ & 12.0 & 11.8 & 12.0\\\hline \end{tabular} \tablefoot{CU Vir (this work), HR 5907 \citep{leto18}, HR 7355 \citep{leto17}.} \end{center} \end{table} \subsection{Purely thermal vs. thermal plus non-thermal X-rays} In Ap/Bp stars the stellar magnetosphere offers several mechanisms capable of producing X-ray emission. The strongly magnetic star CU Vir possesses a primarily dipolar-like non-axisymmetric magnetic structure. Thermal X-ray plasma will be naturally created in wind shocks via the MCWS mechanism, but the weaker winds in late-B\,/\,early-A stars with terminal velocities of $V_{\inf}\approx 600$~km\,s$^{-1}$ are expected to create plasma with post-shock temperatures of a few up to about 10~MK \citep[e.g.][]{dou14}. The 'classical' MCWS plasma is thus a suitable candidate for the cooler plasma component(s) in our spectral models, but insufficient to explain the hot, dominant component around 30~MK. Overall, the main characteristics of CU~Vir derived here, i.e. faint but hard X-ray emission, are virtually the exact opposite of the original motivation for the development of the MCWS model by \cite{bab97} to explain the ROSAT data of the A0p star IQ~Aur, i.e. bright but soft X-ray emission. Clearly, a mechanism that, depending on the adopted model, is producing the hottest plasma or non-thermal component is needed. Furthermore, this component is not a small extra, but actually the dominant contribution to the observed X-ray emission; comparing results from 2011 ({\it XMM-Newton}) with those from 2017 ({\it Chandra}) it is likely also the more variable component. The very hot plasma could arise from activity-like phenomena associated with magnetic spots or plasma captured in the magnetosphere, where in addition re-connection events might occur in the disk-like structure e.g. during break-out or infall events. However, compared to similar Ap/Bp stars like the well studied IQ~Aur \citep{rob11}, the required plasma temperatures of CU~Vir are extreme. We find quasi-quiescent X-rays for IQ~Aur with $\log L_{\rm X} \approx 29.6$~erg\,s$^{-1}$ and $T_{\rm X} \approx 8$~MK and for CU~Vir with $\log L_{\rm X} \approx 28.4$~erg\,s$^{-1}$ and $T_{\rm X} \approx 25$~MK. Similarly hard X-ray spectra are only seen during a flare event in IQ~Aur, which is associated with very hot thermal plasma at several tens of MK as deduced from the detection of a strong 6.7~keV \ion{Fe}{XXV} emission line complex. However, in IQ~Aur this is a transient phenomenon, whereas the very hot component is clearly associated with the quasi-quiescent state of CU~Vir. An alternative scenario, involving a non-thermal component that describes auroral X-ray emission, was proposed by \cite{leto17}, put forward to explain the X-ray emission from fast rotating magnetic early B star HR~7355. It suggests, that besides the thermal emission originating from the MCWS that heats plasma up to a few MK, non-thermal emission should be present in fast rotating magnetic Ap/Bp stars. In this scenario the X-ray emission originates from the non-thermal electrons upon impact on the stellar surface. The electron population would be identical to the one, that is also responsible for the gyrosynchrotron radio emission. The non-thermal electron population able to produce the incoherent gyrosynchrotron emission of CU~Vir has a power-law energy distribution \citep{leto06} with a low energy cutoff of $\approx 100$ keV \citep{tri04}. Such non-thermal electrons, accelerated by magnetic reconnection events occurring in the current sheets regions located far from the star about 15 stellar radii, precipitate toward the stellar surface. The energy budget of these precipitating electrons is compatible with what observed during the Sun flares, where also hard X-ray emission is detected at the footprints of the magnetic loops \citep[see][and references therein]{asch02}. In this case the hard X-ray component would be a power-law component generated by thick target bremsstrahlung emission from the non-thermal electrons. Indeed, CU~Vir could be understood as a down-scaled version of HR~7355, whose radio luminosity as well as its X-ray luminosity is about a factor 30 higher. What causes these phenomena in CU~Vir and similar objects remains so far open, but fast rotation and strong dipolar fields have been identified as common attributes. Further hard X-ray observations at photon energies higher than 5~keV could be very useful to definitively confirm this scenario. The X-ray spectrum of CU~Vir alone is not able to definitely support the scenario where the hard X-ray component has auroral origin. The thermal and the non-thermal scenario are both viable and even both may be at work simultaneously. Furthermore, the thermal component in itself may be a composite that is partly MCWS, magnetospheric activity or reprocessed auroral emission in nature. However, in any case some extraordinary X-ray generation mechanism has to be present in any intrinsic emission scenario. | 18 | 8 | 1808.02367 |
1808 | 1808.07297_arXiv.txt | {The star formation rate (SFR) linearly correlates with the amount of dense gas mass (M$_{dg}$) involved in the formation of stars both for distant galaxies and clouds in our Galaxy. Similarly, the mass accretion rate ($\dot{M}_{\rm acc}$) and the disk mass (M$_{disk}$) of young, Class II stars are also linearly correlated.} {We aim to explore the conditions under which the previous relations could be unified.} {Observational values of SFR, M$_{dg}$, $\dot{M}_{\rm acc}$, and M$_{disk}$ for a representative sample of galaxies, star forming clouds, and young stars have been compiled from the literature. Data were plotted together in order to analyze how the rate of gas transformed into stars and the mass of dense gas directly involved in this transformation relate to each other over vastly different physical systems.} {A statistically significant correlation is found spanning $\sim$ 16 orders of magnitude in each axis, but with large scatter. This probably represents one of the widest ranges of any empirical correlation known, encompassing galaxies that are several kiloparsec in size, parsec-size star-forming clouds within our Galaxy, down to young, pre-main sequence stars with astronomical unit-size protoplanetary disks. Assuming that this global correlation has an underlying physical reason, we propose a bottom-up hypothesis suggesting that a relation between $\dot{M}_{\rm acc}$ and the total circumstellar mass surrounding Class 0/I sources (M$_{cs}$; disk+envelope) drives the correlation in clouds that host protostars and galaxies that host clouds. This hypothesis is consistent with the fact that the SFRs derived for clouds over a timescale of 2 Myr can be roughly recovered from the sum of instantaneous accretion rates of the protostars embedded within them, implying that galactic SFRs averaged over $\sim$ 10-100 Myr should be constant over this period too. Moreover, the sum of the circumstellar masses directly participating in the formation of the protostellar population in a cloud likely represents a non-negligible fraction of the dense gas mass within the cloud.} {If the fraction of gas directly participating in the formation of stars is $\sim$ 1-35$\%$ of the dense gas mass associated with star-forming clouds and galaxies, then the global correlation for all scales has a near unity slope and an intercept consistent with the (proto-)stellar accretion timescale, M$_{cs}$/$\dot{M}_{\rm acc}$. Therefore, an additional critical test of our hypothesis is that the $\dot{M}_{\rm acc}$-M$_{disk}$ correlation for Class II stars should also be observed between $\dot{M}_{\rm acc}$ and M$_{cs}$ for Class 0/I sources with similar slope and intercept.} | \label{Sect:Intro} Star formation encompasses a broad range of physical regimes that account for the collapse of gas to form smaller structures: from star-forming galaxies to individual young stars surrounded by accreting envelopes and disks, passing through intermediate scales including giant molecular clouds, clumps, and dense cores. Almost 60 years ago, \citet{Schmidt59} proposed that the star formation rate (SFR) is a power-law function of the total gas density (atomic and molecular), $\sum$$_{SFR}$ = $K$$\sum$$_{gas}^N$ \citep{Kennicutt98}, with $\sum$$_{SFR}$ and $\sum$$_{gas}$ the SFR and the total gas mass (M$_{gas}$) per unit surface area, and $K$ and $N$ the appropriate constants. The Schmidt-Kennicutt law has been observationally confirmed many times when spanning different types of galaxies, revealing a nonlinear correlation with typical values 1 $<$ $N$ $<$ 2 and significant scatter. Using the HCN molecule to trace dense gas mass (M$_{dg}$; n(H$_{2}$) $>$ 10$^4$ cm$^{-3}$), \citet{Gao04} found a roughly linear ($N$ $\sim$ 1) correlation between the spatially integrated SFR and M$_{dg}$ for a wide range of galaxies including normal spirals, luminous, and ultra-luminous infrared galaxies (LIRGs and ULIRGs, respectively). \citet{Wu05} then found that the spatially integrated SFR in distant star-forming clouds in our Galaxy also correlates linearly with the amount of HCN-emitting gas \citep[see also][]{Jackson13, Stephens16}, showing that this correlation also connected smoothly with the one previously found for galaxies. On this basis, \citet{Wu05} suggested that the formation of stars in other galaxies may be understood in terms of the basic units of star formation in the Milky Way (MW), which they identified as the clouds massive enough to fully sample the initial mass function (IMF) by hosting massive young stars. The measurements were extended to nearby molecular clouds with and without massive stars using dust extinction to more accurately trace dense gas masses and direct counts of young stars to more robustly measure the SFRs \citep{Lada10,Lada12}. These measurements confirm that the correlation is linear and spreads approximately ten orders of magnitude in each axis. In summary, strong lines of evidence support the linearity of the SFR-M$_{dg}$ correlation and that this connects individual Galactic clouds and the various types of star-forming galaxies \citep[see also the recent works in][]{Shimajiri17,Tan18}. The previous view somewhat neglects the role of the actual units in star formation, which are the individual, forming stars. This can be partly attributed to the fact that the SFR, defined as the rate at which gas transforms into stars (in units of M$_{\odot}$ yr$^{-1}$), characterizes physical systems that typically include $\sim$ 10-10$^{11}$ objects. However, there is an equivalent parameter for individual stars that also accounts for the gas-to-star transformation by measuring the rate at which gas falls from the circumstellar environment onto the stellar surface: the stellar mass accretion rate ($\dot{M}_{\rm acc}$, in units of M$_{\odot}$ yr$^{-1}$ too). Direct estimates of $\dot{M}_{\rm acc}$ are only available for optically visible pre-main sequence objects. These are young stars in a relatively advanced stage of early stellar evolution \citep[Class II from the classification of][]{Lada87}, when the optically thick envelope has dissipated and direct observations of the remaining circumstellar disks and the stellar accretion shocks are possible. Their gas disk masses (M$_{disk}$; n(H$_{2}$) $>$ 10$^5$ cm$^{-3}$) are in turn inferred either from observations of different molecules or from submillimeter and millimeter dust continuum emission and an assumed gas-to-dust ratio. Different works indicate that so far the latter method provides more accurate measurements of M$_{disk}$ \citep{Manara16,Rosotti17}. In fact, $\dot{M}_{\rm acc}$ and dust-based M$_{disk}$ measurements are also linearly correlated for large samples of optically visible young stars \citep{Najita07,Mendi12,Najita15}, a result that has been unambiguously confirmed from recent Atacama Large Millimeter/submillimeter Array (ALMA) data at least for the star-forming regions Lupus and Cha I, despite the large scatter \citep{Manara16,Mulders17}. Although one could be tempted to interpret the $\dot{M}_{\rm acc}$-M$_{disk}$ correlation of stars only in terms of small-scale physical processes in the circumstellar disks, this correlation resembles a scaled-down version of the above-mentioned linear correlation linking local clouds and other galaxies. This paper shows that the linear correlation that links galaxies and star-forming clouds within our Galaxy (through SFR and M$_{dg}$) can indeed be connected down to individual, forming stars when physically equivalent measurements ($\dot{M}_{\rm acc}$ and M$_{disk}$) are used for the latter objects. We propose that such global correlation has an underlying explanation that may ultimately contribute to a unified view of star formation across an enormous range of astronomical scale. Section \ref{Sect:stars-clouds} shows that the correlations for young stars and clouds in the MW can be connected. We discuss the origin of such a single correlation under the hypothesis that the relation observed in clouds is driven by the individual (mainly Class 0/I) stars that form them. Section \ref{Sect:global} extends the previous finding and discussion, including data for other galaxies. Finally, Sect. \ref{Sect:conclusions} summarizes the main conclusions. | \label{Sect:conclusions} The gas-to-star transformation rate and the dense gas mass available for star formation show a statistically significant correlation that has been extended from galaxies and Galactic clouds to individual young stars. The correlation spreads $\sim$ 16 orders of magnitude, covering spatial scales from distant (hundreds of Mpc) galaxies with sizes of a few Kpc, to nearby ($\sim$ 150 $pc$) protoplanetary disks with sizes of hundreds of au. Given the difference between the intercepts of the individual linear fits determined for protoplanetary disks, star-forming clouds, and galaxies, it is not clear how physically meaningful the global correlation is. Here we have proposed and explored one possible unifying hypothesis that could explain the global correlation. Our bottom-up hypothesis is that the correlation between the mass accretion rate and the total circumstellar gas mass surrounding young -- mainly Class 0/I -- stars is the underlying cause of the global correlation on all scales. Two main lines of evidence supporting this scenario have been provided. First, the SFRs of molecular clouds are recovered from the sum of empirically based stellar accretion rates. Second, the sum of the individual circumstellar gas masses appears to represent a relevant fraction of the total dense gas mass within such clouds. If that fraction is between $\sim$ 1\%\ and 35 $\%$ of the dense gas mass in clouds and galaxies, then the global correlation has a near unity slope and intercept consistent with the stellar accretion timescale, M$_{cs}$/$\dot{M}_{\rm acc}$. Indeed, a crucial test of our hypothesis is that the $\dot{M}_{\rm acc}$-M$_{disk}$ linear correlation on astronomical unit scales for Class II stars should also be observed between $\dot{M}_{\rm acc}$ and M$_{cs}$ on (sub-)parsec scales for Class 0/I sources. Our hypothesis also explains that the scatter in the correlation decreases from stars to galaxies, and implies that the average SFR of galaxies remains constant on a timescale of $\sim$ 10-100 Myr. The $\dot{M}_{\rm acc}$-M$_{disk}$ and SFR-M$_{dg}$ correlations in stars, clouds, and galaxies have been independently studied from physical contexts relatively isolated from each other. Our result suggests a global approach and that theoretical efforts should consider all scales and physical systems involved in a single correlation. | 18 | 8 | 1808.07297 |
1808 | 1808.09243_arXiv.txt | {Very long baseline interferometry (VLBI) observations at 86$\,$GHz (wavelength, $\lambda = 3$\,mm) reach a resolution of about 50 $\mu$as, probing the collimation and acceleration regions of relativistic outflows in active galactic nuclei (AGN). The physical conditions in these regions can be studied by performing 86\,GHz VLBI surveys of representative samples of compact extragalactic radio sources.} {To extend the statistical studies of compact extragalactic jets, a large global 86\,GHz VLBI survey of 162 compact radio sources was conducted in 2010--2011 using the Global Millimeter VLBI Array (GMVA).} { The survey observations were made in a snapshot mode, with up to five scans per target spread over a range of hour angles in order to optimize the visibility coverage. The survey data attained a typical baseline sensitivity of 0.1 Jy and a typical image sensitivity of 5 mJy/beam, providing successful detections and images for all of the survey targets. For 138 objects, the survey provides the first ever VLBI images made at 86 GHz. Gaussian model fitting of the visibility data was applied to represent the structure of the observed sources and to estimate the flux densities and sizes of distinct emitting regions (components) in their jets. These estimates were used for calculating the brightness temperature ($T_\mathrm{b}$) at the jet base (core) and in one or more moving regions (jet components) downstream from the core. These model-fit-based estimates of $T_\mathrm{b}$ were compared to the estimates of brightness temperature limits made directly from the visibility data, demonstrating a good agreement between the two methods.} { The apparent brightness temperature estimates for the jet cores in our sample range from $2.5 \times 10^{9}$\,K to $ 1.3\times 10^{12}$\,K, with the mean value of $1.8 \times 10^{11} $\,K. The apparent brightness temperature estimates for the inner jet components in our sample range from $7.0 \times 10^{7}$\,K to $4.0 \times 10^{11}$\,K. A simple population model with a single intrinsic value of brightness temperature, $T_\mathrm{0}$, is applied to reproduce the observed distribution. It yields $T_\mathrm{0} = (3.77^{+0.10}_{-0.14}) \times 10^{11}$\,K for the jet cores, implying that the inverse Compton losses dominate the emission. In the nearest jet components, $T_\mathrm{0} = (1.42^{+0.16}_{-0.19}) \times 10^{11}$\,K is found, which is slightly higher than the equipartition limit of $\sim5\times 10^{10}$\,K expected for these jet regions. For objects with sufficient structural detail detected, the adiabatic energy losses are shown to dominate the observed changes of brightness temperature along the jet.} {} | VLBI (Very long baseline interferometry) observations at 86\,GHz (wavelength, $\lambda=3$\,mm) enable detailed studies to be made of compact radio sources at a resolution of $\sim $ (40 --100) $\mu$as. This resolution corresponds to linear scales as small as 10$^3$--10$^4$ Schwarzschild radii and uncovers the structure of the jet regions where acceleration and collimation of the flow takes place \citep{vlahakis2004,lee2008,lee2016,asada2014,boccardi2016,mertens2016}. To date, five 86\,GHz VLBI surveys have been conducted \citep[][see Table~\ref{tb:surveys}]{beasley1997,lonsdale1998,rantakyro1998,lobanov2000,lee2008}, with the total number of objects imaged reaching just over a hundred. No complete sample of objects imaged at 86\,GHz has been established so far. Recent works \citep[{e.g.},][]{homan2006,cohen2007,lister2016} have demonstrated that high-resolution studies of complete (or nearly complete) samples of compact jets yield a wealth of information about the intrinsic properties of compact extragalactic flows. Measuring brightness temperature in a statistically viable sample enables the performance of detailed investigations of the physical conditions in this region. The distribution of observed brightness temperatures, $T_\mathrm{ b}$, derived at 86\,GHz can be combined with the $T_\mathrm{ b}$ distributions measured at lower frequencies \citep[{e.g.},][]{kovalev2005}. This can help to constrain the bulk Lorentz factor, $\Gamma_\mathrm{j}$, and the intrinsic brightness temperature, $T_\mathrm{0}$, of the jet plasma, using different types of population models of relativistic jets \citep{vermeulen1994,lobanov2000,lister2003,homan2006}. Changes of $T_\mathrm{0}$ in the compact jets with frequency can be used to distinguish between the emission coming from accelerating or decelerating plasma and from electron-positron or electron-proton plasma. Theoretical models predict $T_\mathrm{0} \propto \nu^{\epsilon}$, with $\epsilon \approx 2.8$, below a critical frequency $\nu_\mathrm{ break}$ at which energy losses begin to dominate the emission \citep{marscher1995}. Above $\nu_\mathrm{break}$, $\epsilon$ can vary from $-$1 to $+$1, depending on the jet composition and dynamics. By measuring this break and the power-law slopes above and below, it would be possible to distinguish between different physical situations in the compact jets. Previous studies \citep{lobanov2000,lee2008} indicate that the value of $\nu_\mathrm{break}$ is likely to be below 86\,GHz. Indeed, a compilation of brightness temperatures measured at 2, 8, 15, and 86\,GHz \citep{lee2008} indicates that brightness temperatures measured at 86\,GHz are systematically lower and $\nu_\mathrm{break}$ can be as low as 20\,GHz. This needs to be confirmed on a complete sample observed at 86\,GHz. If $T_\mathrm{0}$ starts to decrease at 86$\,$GHz, there will be only a few sources suitable for VLBI $> 230$\,GHz and higher frequencies. Such a decrease of $T_\mathrm{0}$ will also provide a strong argument in favour of the decelerating jet model or particle-cascade models as discussed by \cite{marscher1995}. In view of these arguments, it is important to undertake a dedicated 86$\,$GHz VLBI study of a larger complete sample of extragalactic radio sources. In this paper, we present results from a large global VLBI survey of compact radio sources carried out in 2010--2011 with the Global Millimeter VLBI Array (GMVA)\footnote{http://www3.mpifr-bonn.mpg.de/div/vlbi/globalmm/.}. This survey has provided images of 162 unique radio sources, increasing the total number of sources ever imaged with VLBI at 86$\,$GHz by a factor of 1.5. The combined database resulting from this survey and \cite{lee2008} comprises 256 sources. This information provides a basis for investigations of the collimation and acceleration of relativistic flows and probing the physical conditions in the vicinity of supermassive black holes. The survey data reach a typical baseline sensitivity of 0.1\,Jy and a typical image sensitivity of 5\,mJy/beam. A total of 162 unique compact radio sources have been observed in this survey and all the sources are detected and imaged. With the present survey, the overall sample of compact radio sources imaged with VLBI at 86\,GHz is representative down to $\sim 0.5$\,Jy for J2000 declinations of $\delta\ge 15^{\circ}$. Section~\ref{sc:obs} describes the source selection and the survey observations. In Section~\ref{sc:data}, we describe the data processing, amplitude and phase calibration, imaging, model fitting procedures and a method for estimating errors of the model-fit parameters. Section~\ref{sc:results} describes the images and derived parameters of the target sources. Examples of images of four selected sources obtained from the survey data are presented in Section~\ref{sc:images} (the complete set of images of all of the target sources is presented in Appendix available electronically). Brightness temperatures of the survey sources are derived and discussed in Section~\ref{sc:tb}. Section~\ref{sc:model} describes population modelling of the brightness temperature distribution observed at the base (VLBI core) of the jet and in the innermost moving jet components. Evolution of the observed brightness temperature along the jet is studied in Section~\ref{sc:evol} for the target sources with sufficient extended emission detected. | \label{sc:discussion} \subsection{Modelling the observed brightness temperatures} \label{sc:model} The brightness temperature distribution can be used for obtaining estimates of the conditions in the extra galactic radio sources and to test the models proposed for the inner jets \citep{marscher1995,lobanov2000,homan2006,lee2008}. A basic population model \citep{lobanov2000} can be used for representing the observed brightness temperature distribution under the assumption that the jets have the same intrinsic brightness temperature, $T_\mathrm{0}$, Lorentz factor, $\Gamma_\mathrm{j}$, and synchrotron spectral index, $\alpha$ ($S_\nu \propto \nu^\alpha$), and they are randomly oriented in space (within the limits of viewing angles, $\theta$ , required by Doppler boosting bias). The jets are also assumed to remain straight within the spatial scales ($\sim 0.5-10\,\mathrm{pc}$) probed by the observations. The assumptions of single values of $T_\mathrm{0}$ and $\Gamma_\mathrm{j}$ describing the whole sample are clearly simplified, as jets are known to feature a range of Lorentz factors \citep[see][and references therein]{lister2016}. However, as has been shown earlier \citep{lobanov2000}, factoring a distribution of Lorentz factors into the present model is not viable without amending the brightness temperature measurements with additional information, preferably about the apparent speeds of the target sources. We are currently compiling such a combined database, and will engage in a more detailed modelling of the compact jets after the completion of this database. In a population of jets described by the settings summarized above, the measured brightness temperature, $T_\mathrm{b}$, is determined solely by the relativistic Doppler boosting of the jet emission. Therefore, the observed brightness temperature , $T_\mathrm{b}$ , can be related to the intrinsic brightness temperature, $T_0$, so that $T_\mathrm{b} = T_0\,\delta^{1/\epsilon}$, where the power index $\epsilon$ is $1/(2-\alpha)$ for a continuous jet (steady state jet) and $1/(3-\alpha)$ for a jet with spherical blobs (or optically thin ``plasmoids''), and $\delta$ is the Doppler factor. The probability of finding a radio source with the brightness temperature $T_\mathrm{b}$ in such a population of sources is \begin{equation} \label{eqn:prob_equation} p(T_\mathrm{b}) \propto \Bigg[ \frac{2\,\Gamma_{j}\left(T_\mathrm{b}/{T_\mathrm{0}}\right)^\epsilon - \left(T_\mathrm{b}/{T_\mathrm{0}}\right)^{2\epsilon} - 1} {\Gamma_\mathrm{j}^{2} - 1}\Bigg]^{\frac{1}{2}}\,. \end{equation} \begin{figure}[ht!] \centering \includegraphics[width=0.52\textwidth,trim={1cm 0 1cm 1.3cm},clip=true]{Plots/population_model_cores.eps} \caption{Distribution of the brightness temperatures, $T_\mathrm{b}$ , measured in the core components and represented by the population models calculated for $\Gamma_\mathrm{j}=10$ and different values of $T_\mathrm{0}$. The best approximation of the observed $T_\mathrm{b}$ distribution is obtained with $T_\mathrm{o,core}$ = $(3.77^{+0.10}_{-0.14}) \times 10^{11}$\,K. For better viewing of the observed distribution, one core component with a very high $T_\mathrm{b}$ = $ 5.5 \times 10^{12}$\,K for the source BL Lac obtained from \cite{lee2008} is not shown but is included in the modelling.} \label{fg:tbcores_all} \end{figure} \begin{figure}[ht!] \centering \includegraphics[width=0.52\textwidth,trim={1cm 0 1cm 1.3cm},clip=true]{Plots/population_model_jets.eps} \caption{Distribution of the brightness temperatures, $T_\mathrm{b}$ , measured in the inner jet components and represented by the population models calculated for $\Gamma_\mathrm{j}=10$ and different values of $T_\mathrm{0}$. The best approximation of the observed $T_\mathrm{b}$ distribution is obtained with $T_\mathrm{o,jet} = (1.42^{+0.16}_{-0.19}) \times 10^{11}$\,K. } \label{fg:tbjetcomponents_all} \end{figure} The lower end of the observed distribution of brightness temperatures depends on the sensitivity of VLBI data since the flux of the observed sample is biased by Doppler boosting \citep{lobanov2000}. The lowest brightness temperature that can be measured from our data, $T_\mathrm{b,sens}$ , can be obtained from \begin{equation} \label{eqn:Tb_sens} T_{\mathrm{b,sens}} \mathrm{[K]} = 1.65 \times 10^{5} \left(\frac{\sigma_{\mathrm{rms}}}{\mathrm{mJy/beam}}\right) \left(\frac{b}{\mathrm{mas}}\right)^{-2}\,, \end{equation} \noindent where $\sigma_\mathrm{rms}$ is the array sensitivity in mJy/beam and $b$ is the average size of the resolving beam. In this survey, the typical observation time on a target source $\Delta t$ is 20 minutes and bandwidth is 128\,MHz. Therefore, the value of beam size for the sources in this survey is 0.12 mas and the $\sigma_\mathrm{rms}$ of the array is 0.54 mJy/beam. Thus, we have obtained a 3$\sigma$ level estimation of $T_\mathrm{b,sens}$ as 2.0 $\times$ 10$^{8}$ K using Equation~\ref{eqn:Tb_sens}, which is set as the lowest brightness temperature in modelling. We normalize the results obtained from Equation (\ref{eqn:prob_equation}) to the number of objects in the lowest bin of the histogram. For our modelling, we first make a generic assumption of $\alpha=-0.7$ (homogeneous synchrotron source) and use $\Gamma_\mathrm{j} \approx 10$ implied from the kinematic analysis of the MOJAVE VLBI survey data \citep{lister2016}. We assume that the jet is continuous, so $\epsilon =$ 0.37 is taken. For the population modelling analysis, we have included the data from \cite{lobanov2000}, \cite{lee2008}, and the present survey, yielding a final database of 271 VLBI core components and 344 jet components. For objects with multiple measurements, we have used the median of the measurements. The resulting model distributions obtained for various values of $T_\mathrm{0}$ are shown in Figures~(\ref{fg:tbcores_all}--\ref{fg:tbjetcomponents_all}) for the VLBI cores and the inner jet components, respectively. This approach yields $T_\mathrm{0,core}$ = $(3.77^{+0.10}_{-0.14}) \times 10^{11}$\,K for the VLBI cores and $T_\mathrm{0,jet} = (1.42^{+0.16}_{-0.19}) \times 10^{11}$\,K for the inner jet components. The estimated $T_\mathrm{0,core}$ is in good agreement with the inverse Compton limit of $\simeq$ 5.0 $\times$ 10$^{11}$\,K \citep{kellermann1969}, beyond which the inverse Compton effect causes rapid electron energy losses and extinguishes the synchrotron radiation. The inferred $T_\mathrm{0,jet}$ of jet components are about a factor of three higher than the equipartition limit of $\simeq$ 5 $\times$ 10$^{10}$\,K \citep{readhead1983} for which the magnetic field energy and particle energy are in equilibrium. This may indicate that opacity is still non-negligible in these regions of the flow. The intrinsic brightness temperature obtained for the cores is within the upper limit $5.0 \times 10^{11}$\,K predicted for the population modelling of the cores \citep{lobanov2000}. A simultaneous fit for $T_\mathrm{0}$ and $\Gamma_\mathrm{j}$ is impeded by the implicit correlation, $T_\mathrm{0} \propto \Gamma_\mathrm{j}^{a}$ (with $a \approx 2$--3), between these two parameters, as implied by Equation (\ref{eqn:prob_equation}). This is also illustrated in Figure~\ref{fg:simfit}, from which a dependence $T_\mathrm{0}[\mathrm{K}] \approx (7.7\times 10^{8})\,\Gamma_\mathrm{j}^{2.7}$ can be inferred for the fit to the brightness temperatures measured in the VLBI cores. This correlation between $T_\mathrm{0}$ and $\Gamma_\mathrm{j}$ precludes simultaneously fitting for both these parameters, and hence the \begin{figure*}[ht!] \centering \includegraphics[width=0.90\textwidth,clip=true]{Plots/population_model_fit.eps} \caption{Two-dimensional ${\chi}^2$ distribution plot in the $\Gamma_\mathrm{j}$ -- $T_\mathrm{0}$ space, calculated for the brightness temperatures measured in the VLBI cores. The blank area shows the ranges of the parameter space disallowed by the observed distribution. The distribution of the ${\chi}^2$ values indicates a ($\Gamma_\mathrm{j}$--$T_\mathrm{0}$) correlation, with $T_\mathrm{0}[\mathrm{K}]\approx (7.7\times 10^{8})\,\Gamma_\mathrm{j}^{2.7}$, thus precluding a simultaneous fit for $\Gamma_\mathrm{j}$ and $T_\mathrm{0}$.} \label{fg:simfit} \end{figure*} \noindent Lorentz factor has to be constrained (or assumed) separately. One should also keep in mind that this correlation results from the model description and does not have an immediate physical implication. Equation (\ref{eqn:prob_equation}) clearly shows that the predicted distribution of $T_\mathrm{b}$ is valid within the range \begin{equation} \label{eqn:prob_range} (\Gamma_\mathrm{j} - \sqrt{ \Gamma_\mathrm{j}^{2} - 1}) \leq (\frac{T_\mathrm{0}}{T_\mathrm{b}})^{\epsilon} \leq (\Gamma_\mathrm{j} + \sqrt{ \Gamma_\mathrm{j}^{2} - 1})\,. \end{equation} The region outside this range is represented by the blank area in Figure~\ref{fg:simfit}. The intrinsic brightness temperature we obtained is higher than the mean and median observed brightness temperature $T_\mathrm{b}$. This is readily explained by the Doppler deboosting. For a given viewing angle, $\theta$, sources with $\Gamma_\mathrm{j}>1/\theta$ would be deboosted so that the observed brightness temperature will be reduced below its intrinsic value. It can be easily shown that the observed and intrinsic brightness temperatures are equal if the jet viewing angle is given by \begin{equation} \label{eqn:viewing angle} \theta_\mathrm{eq} = \arccos \left[ \frac{1- (1/\Gamma_\mathrm{j})\,\,(T_\mathrm{0}/T_\mathrm{b})^{\epsilon}} {\sqrt{1-\Gamma_\mathrm{j}^{-2}}} \right]\,. \end{equation} For the VLBI cores, the mean of the observed $T_\mathrm{b}$ is $1.8 \times 10^{11}$\,K and intrinsic $T_\mathrm{0,core}$ is $3.77 \times 10^{11}$\,K, therefore, the resulting $\theta_\mathrm{eq}$ = 29$^{\circ}$ for $\Gamma_\mathrm{j}$ = 10 and $\epsilon$ = 0.37. In this case any object observed at a larger viewing angle would be deboosted resulting in a lower observed $T_\mathrm{b}$ than intrinsic $T_\mathrm{0}$. \subsection{Testing the adiabatic expansion of jets} \label{sc:evol} As discussed in Sections~\ref{sc:tb} and \ref{sc:model}, intrinsic $T_\mathrm{0}$ and the observed $T_\mathrm{b}$ in core and jets show that the brightness temperature drops by approximately a factor of two to ten already on sub-parsec scales in the jets. This evolution might occur with the inverse Compton, synchrotron, and adiabatic losses subsequently \begin{figure*}[ht!] \centerline{\includegraphics[width=0.5\textwidth]{Plots/adiabatic_J2322+5057.eps} \includegraphics[width=0.5\textwidth]{Plots/adiabatic_J1033+6051.eps} } \centerline{\includegraphics[width=0.5\textwidth]{Plots/adiabatic_3C345.eps} \includegraphics[width=0.5\textwidth]{Plots/adiabatic_0716+714.eps} } \caption{Changes of the brightness temperature as a funcion of jet width for four sources -- J2322+5037, J1033+6051, 3C 345, and 0716+714 from this survey. Blue squares denote the measured $T_\mathrm{b}$ from this survey. The red circles connected with a dotted line represent theoretically expected $T_\mathrm{b}$ under the assumption of adiabatic jet expansion. The initial brightness temperature in each jet is assumed to be the same as that measured in the VLBI core.} \label{fig:adiabatic} \end{figure*} dominating the energy losses \citep[cf.,][]{marscher1995,lobanov1999}. For four objects in our data (3C84, 0716+714, 3C454.3, and J2322+507) for which multiple jet components have been identified during the model fitting, it is possible to use the brightness temperatures of the jet components to test whether the evolution of the jet brightness on sub parsec scales could be explained by adiabatic energy losses \citep{marscher1985}. For this analysis, we assume that the jet components are independent relativistic shocks embedded in the jet plasma, which has a power-law distribution $ N(E)\,\mathrm{dE} \propto E^{-s}\mathrm{dE,}$ where $s$ is the energy spectral index that depends on spectral index $\alpha$ as $\alpha = (1-s)/2 $, and is pervaded by the magnetic field $B \propto$ $d^{-a}$, where $d$ is the width of the jet and $a$ depends on the type of magnetic field ($a =$ 1 for poloidal magnetic field and 2 for toroidal magnetic field). With these assumptions, we can relate the brightness temperatures, $T_\mathrm{b,J}$, of the jet components to the brightness temperature, $T_\mathrm{b,C}$, of the core \citep{lobanov2000,lee2008}, \begin{equation} \label{eqn:adiab} T_\mathrm{b,J} = T_\mathrm{b,C} \left(\frac{d_\mathrm{J}}{d_\mathrm{C}}\right)^{-\xi}\,, \end{equation} where $d_\mathrm{J}$ and $d_\mathrm{C}$ are the measured sizes of the jet component and core, respectively, and \begin{equation} \xi = \frac {2(2s+1)+3a(s+1)}{6}\,. \end{equation} Assuming the synchrotron emission with spectral index $\alpha$ = -0.5, we use $s = 2.0$ and adopt $a = 1.0$ for the description of the magnetic field in the jet. With these assumptions, we calculate the predicted $T_\mathrm{b,J}$ for individual jet components and compare them in Figure~\ref{fig:adiabatic} to the measured brightness temperatures. The measured and predicted values of brightness temperature agree well, and this suggests that the jet components can be viewed as adiabatically expanding relativistic shocks \citep[cf.,][]{kadler2004,pushkarev2012,kravchenko2016}. | 18 | 8 | 1808.09243 |
1808 | 1808.07068_arXiv.txt | Young planets offer a direct view of the formation and evolution processes that produced the diverse population of mature exoplanet systems known today. The repurposed \emph{Kepler} mission \emph{K2} is providing the first sample of young transiting planets by observing populations of stars in nearby, young clusters or stellar associations. We report the detection and confirmation of two planets transiting \targtight, an M2.5 dwarf in the 650\,Myr old Praesepe open cluster. Using our notch-filter search method on the \emph{K2} lightcurve, we identify planets with periods of 5.84\,d and 19.66\,d. This is currently the second known multi-transit system in open clusters younger than 1\,Gyr. The inner planet has a radius of 2.27$_{-0.16}^{+0.20}$\,R$_\oplus$ and the outer planet has a radius of 2.77$_{-0.18}^{+0.20}$\,R$_\oplus$. Both planets are likely mini-Neptunes. These planets are expected to produce radial velocity signals of 3.4 and 2.7\,m/s respectively, which is smaller than the expected stellar variability in the optical ($\simeq$30\,m/s), making mass measurements unlikely in the optical, but possible with future near-infrared spectrographs. We use an injection-recovery test to place robust limits on additional planets in the system, and find that planets larger than 2\,R$_\oplus$ with periods of 1-20\,d are unlikely. | Planets and their host stars can change dramatically over their lifetimes. Their structural, orbital, and atmospheric properties are all expected to evolve through interactions with their host star \citep[e.g.,][]{Ehrenreich2015}, the protoplanetary disk from which they formed \citep[e.g.,][]{Cloutier2013}, other planets in the system \citep[e.g.,][]{Chatterjee2008}, and the greater stellar environment \citep[e.g.,][]{Cai2018}. Understanding the underlying drivers and relative importance of these evolutionary mechanisms is critical for revealing the early sculpting of planetary systems and the conditions that give rise to the diversity of mature planetary systems revealed by {\it Kepler} and earlier exoplanet surveys \citep[e.g.,][]{2014Natur.513..336L,2016ApJ...817L..13I}. Exoplanets likely evolve the most during their first Gyr \citep[e.g.,][]{Adams2006, Mann2010,Lopez2012}, and planets $<$1\,Gyr old are therefore powerful probes of the important drivers of exoplanet evolution. Fortuitously, the repurposed {\it Kepler} mission, {\it K2} \citep{Howell2014}, has surveyed a number of young clusters and star forming regions spanning $<10$Myr \citep[Taurus-Auriga;][]{Kraus2017a}, to $\simeq$650\,Myr \citep[Praesepe and Hyades;][]{Martin2018}, with Upper Scorpius \citep[$\simeq$10\,Myr;][]{Pecaut2012} and the Pleiades \citep[$\simeq$112\,Myr;][]{Dahm2015} spanning intermediate ages. The Zodiacal Exoplanets in Time (ZEIT) survey \citep{zeit1} was designed to identify and characterize transiting planets in these young clusters and star forming regions using the precise photometry from {\it K2} \citep{VanCleve2016}. The greater goal is to better understand how planets form and evolve by comparing the statistical properties of exoplanets of different ages and to older systems found during the original {\it Kepler} mission \citep{Borucki2010,Thompson2018}. Thus far we have identified planets in Hyades \citep{Vanderburg2018}, Praesepe \citep{zeit4}, and Upper Scorpius \citep{zeit3}, many of which were found near-simultaneously by similar surveys focusing on exoplanets in young stellar associations \citep[e.g.,][]{Obermeier2016, David2016b, Pepper2017, Ciardi2018, Livingston2018}. Multi-transiting planetary systems are uniquely useful for studying stellar and planetary properties. In cases where the planets' eccentricities can be independently constrained \citep[e.g., through dynamics;][]{Deck2016,Gillon2017}, multi-transiting systems can be used to constrain stellar densities with a precision that rivals eclipsing binaries \citep[e.g.,][]{mann17a}. Even with no information on the host star properties, differences between the measured transit duration of planets in the same system can be used to measure the relative eccentricities \citep{Kipping2012}. Multi-transit systems where planets undergo transit timing variations offer the best opportunity to measure the masses of small planets \citep[e.g.,][]{Deck2015,Hadden2017}. Lastly, these systems provide a measurement of the mutual inclination of planets, a probe of the entropy of a system and hence the role of dynamical disruptions from their (expected) initially flat configuration \citep[e.g.,][]{Figueira2012,Ballard2016}. Multi-transiting systems in clusters offer a unique route to study the dynamical properties of planets with known (young) ages. So far, there is only one known multi-transiting system in an open cluster: K2-136, a three-planet system in the 650\,Myr old Hyades cluster \citep{zeit6, Livingston2018,Ciardi2018}. Here, we present the discovery of two planets transiting the 650\,Myr old Praesepe cluster star \targ (JS 597; \citealt{jones91}) from its {\it K2} light curve. \targtight\ hosts two super-Earth/mini-Neptune-sized planets in short ($\approx$6 and $\approx$20 day) period orbits. We describe our discovery and follow-up observations in Section \ref{observationsanddatareduction}, and we describe our analysis to determine stellar properties in Section \ref{sec:stelpars}. In Section \ref{limitsonadditionalplanets} we place limits on additional planets in the system, and in Sections \ref{transitfitting} and \ref{fpp} we describe our transit fitting to determine stellar parameters and our false positive probability analysis. Finally, in Section \ref{discussion}, we discuss the implications from discovering a second multi-transiting cluster system. | \label{discussion} We have reported the discovery and characterization of a two planet transiting system in the Praesepe open cluster. There are now several detected transiting planets in young open clusters and associations observed by \emph{K2}, though \targ is one of only two multiple-planet systems, the other being K2-136, a three transiting-planet system in the Hyades open cluster \citep{zeit6}. \targtight\,b and c are both likely mini-Neptunes, and both sit near the upper envelope of the field mass--panet radius distribution, as is seen for other planets in intermediate-age clusters. These two planets continue the trend of young open clusters M dwarfs hosting planets of larger radii than have been observed for planets transiting older field population dwarfs from the original \emph{Kepler} sample \citep{zeit4,dressing15}. Figure \ref{fig:popradii} shows the planet radii and host star masses of the M-dwarf hosted young planets identified in the ZEIT survey \citep{zeit1,zeit3,zeit4,zeit6}, including \targtight\,b/c, compared to older transiting systems. The possible inflation in radii at $\sim$650\,Myr may be a sign of ongoing atmosphere loss (e.g., \citealt{Lopez2012}). With further completeness testing on the entire sample of Hyades and Praesepe stars observed by \emph{K2} a measure of the rate and significance of the potential radius difference could be measured. \begin{figure} \centering \includegraphics[width=0.49\textwidth]{Praesepe_size_single-eps-converted-to.pdf} \caption{Host star mass and planet radii for the seven transiting planetary systems in Praesepe and the Hyades from \emph{K2} C4/5 \citep{zeit1,zeit4,zeit6} and those presented in this paper from C16, compared to older M-dwarf hosted planets from the original \emph{Kepler} samples \citep{dressing15}. The 650\,Myr Praesepe and Hyades planet population have larger radii than those hosted by older M dwarfs. The single 10\,Myr old planet in Upper Scorpius (K2-33\,b; \citealt{zeit3}) is also significantly larger than the older planets.} \label{fig:popradii} \end{figure} Systems with multiple transiting planets offer the potential for many science cases not possible with single transit systems. In particular, eccentricity and stellar density can strongly constraint each other \citep{vaneylen15,mann17a}. Planet masses for multiple systems can also be measured from transit timing variations (TTV's) \citep{hadden17}. Though we did not explicitly test for TTV's, the detection of TTV's in the \emph{K2} dataset is unlikely; similar size planets show variations of $<$15\,min, which is smaller than the long-cadence timing of $\simeq$30\,mins. In particular, TTV's on planet b due to planet c are expected to be very small ($<$1\,min) given that the orbital periods are very far from a resonance \citep{agol05}. One scenario where TTV detection could be possible involves the presence of a third planet in or near e.g., a 2:1 resonance with the inner planet b. Such a planet would have to be approximately earth-mass to have avoided detection in the \emph{K2} lightcurve. The TTV amplitude from such a planet on the ephemeris of \targtight\,b, assuming zero eccentricities, is 5-15\,min depending on the proximity to resonance \citep{agol05}. The currently available long-cadence data from \emph{K2} is particularly unsuited to the science cases described above. However, \targ is highly amenable to follow-up photometry. Both planets are large enough that ground-based facilities could resolve their transits ($\simeq$3\,mmag), though the faintness of the host star ($r\simeq 16$mag) may be prohibitive for small apertures at high cadence. Shorter cadence data resolving ingress and egress shapes can place stronger constraints on eccentricity, and offer suggestions as to the types of formation mechanisms responsible for forming these two short-period planets. Space-based follow-up with the Hubble Space Telescope or Spitzer is possible for both planets. In Spitzer channel 1 ($\simeq$3.5\,$\mu m$; \citealt{hora08}) \targ is $\simeq$12\,mag \citep{wise10} and in a 2\,min exposure a SNR of 500\,pmm is possible. This is sufficient to resolve the transit shape from even a single transit. Follow-up spectroscopy to measure the masses of \targtight\,b,c may not be possible. \targ shows stellar variability with a period 22.8\,days and photometric amplitude of $\simeq$3\%. If the star is seen equator-on, this amplitude of variability is expected to produce RV variability of $\simeq$30\,m/s in a similar band as \emph{K2}. Using the mass-radius relation for planets from \citet{weiss14} and the radii inferred from our transit fitting, we find that \targtight\,b,c have likely masses of 5.8\,M$_\oplus$ and 7\,M$_\oplus$ respectively. Assuming circular orbits and the stellar properties derived above, these masses correspond to radial velocity semi-amplitudes of 3.4\,m/s and 2.7\,m/s respectively. The amplitude of these signals is significantly smaller than the expected stellar rotations signal. Moving to the near-infrared, where the stellar variability is expected to have significantly smaller amplitude, could alleviate this problem in combination with our prior knowledge of the rotation period of the star. | 18 | 8 | 1808.07068 |
1808 | 1808.04541_arXiv.txt | {Eruptions from coronal bright points (CBPs) are investigated in a two part study.} {The present study aims to explore in full detail the morphological and dynamical evolution of these eruptions in the context of the full lifetime evolution of CBPs. A follow-up study employs data-driven modelling based on a relaxation code to reproduce the time evolution of the magnetic field of these eruptive CBPs, and provide an insight on the possible causes for destabilisation and eruption.} {Observations of the full lifetime of CBPs in data taken with the Atmospheric Imaging Assembly (AIA) on board the Solar Dynamics Observatory in four passbands, He~{\sc ii}~304~\AA, Fe~{\sc ix/x} 171~\AA, Fe~{\sc xii}~193~\AA, and Fe~{\sc xviii}~94~\AA\ are investigated for the occurrence of plasma ejections, micro-flaring, mini-filament eruptions and mini coronal mass ejections (mini-CMEs). Data from the Helioseismic Magnetic Imager are analysed to study the longitudinal photospheric magnetic field evolution associated with the CBPs and related eruptions.} {First and foremost, our study shows that the majority (76\%) of quiet Sun CBPs (31 out of 42 CBPs) produce at least one eruption during their lifetime. From 21 eruptions in 11 CBPs, 18 occur in average $\sim$17~hrs after the CBP formation for an average lifetime of the CBPs in AIA~193~\AA\ of $\sim$21~hrs. This time delay in the eruption occurrence coincides in each BP with the convergence and cancellation phase of the CBP bipole evolution during which the CBPs become smaller until they fully disappear. The remaining three happen 4 -- 6~hrs after the CBP formation. In sixteen out of 21 eruptions the magnetic convergence and cancellation involve the CBP main bipoles, while in three eruptions one of the BP magnetic fragments and a pre-existing fragment of opposite polarity converge and cancel. In one BP with two eruptions cancellation was not observed. The CBP eruptions involve in most cases the expulsion of chromospheric material either as elongated filamentary structure (mini-filament, MF) or a volume of cool material (cool plasma cloud, CPC), together with the CBP or higher overlying hot loops. Coronal waves were identified during three eruptions. A micro-flaring is observed beneath all erupting MFs/CPCs. It remains uncertain whether the destabilised MF causes the micro-flaring or the destabilisation and eruption of the MF is triggered by reconnection beneath the filament. In most eruptions, the cool erupting plasma obscures partially or fully the micro-flare until the erupting material moves away from the CBP. From 21 eruptions 11 are found to produce mini-CMEs. The dimming regions associated with the CMEs are found to be occupied by both the `dark' cool plasma and areas of weakened coronal emission caused by the depleted plasma density. } {The present study demonstrates that the evolution of small-scale loop structures in the quiet Sun determined by their magnetic footpoint motions and/or ambient field topology, evolve into eruptive phase that triggers the ejection of cool and hot plasma in the corona. } | \label{sect:intro} Coronal bright points (CBPs) are one of the most typical small-scale phenomena in the solar corona. The earliest discovery of CBPs was in X-rays which initially led to the introduction of the term ``X-ray BP'' \citep{vaiana1970a}. X-ray BPs were seen in rocket mission observations as point-like structures with typical sizes ranging from 10\arcsec\ to 50\arcsec\ that form a bright core of $\sim$10\arcsec\ \citep{1974ApJ...189L..93G}. Their lifetimes differ when observed in different imaging channels, i.e. when observed at different temperatures. CBPs were found to have a lifetime in X-rays from 2 to 48 hours with an average lifespan of 8 hours \citep{1976SoPh...49...79G}. Later, using data from the Yohkoh X-ray telescope, \citet{1993AdSpR..13...27H} studied the lifetimes of 514 X-ray BPs and established that coronal hole BPs (34\% of the total) have an average existence of 12~hrs 36~min, the QS BPs (60\%) 13~hrs and 11~hrs for active region BPs (6\%). Studies based on extreme-ultraviolet (EUV) observations \citep[e.g.,][]{1981SoPh...69...77H, 2001SoPh..198..347Z} show that the lifetimes of CBPs in EUV range from a few minutes to a few days with an average value of 20~hrs \citep{2001SoPh..198..347Z}. For review of on the observational properties of CBPs and their modelling see Madjarska (in prep). High-resolution spectroheliograms in Fe~{\sc xiv}~284~\AA\ taken with the extreme ultraviolet spectrograph on the Apollo Telescope Mount on board {\it Skylab} first revealed that CBPs consist of small-scale fast evolving loops \citep{1979SoPh...63..119S} which was later confirmed by the high-resolution TRACE data \citep[e.g.,][]{2004A&A...418..313U}. \citet{1990ApJ...352..333H} observed that the peaks of the emission in CBPs in six simultaneously recorded different spectral lines with various formation temperatures are not always co-spatial which made the authors conclude that CBPs are possibly composed of small-scale loops at different temperatures. From near limb spectroheliogram observations of CBPs that cover chromospheric, transition region and coronal lines taken with the Harvard EUV experiment in 1973, \cite{1981SoPh...69...77H} found that the coronal emission in CBPs is located a few arcseconds above the CBP chromospheric emission sources which should be assumed if CBPs are formed by magnetic loops connecting bipolar regions. Recent spectroscopic observations from the Extreme Ultraviolet Spectrometer (EIS) on board the Hinode satellite further confirmed this by finding that the transition region bipolar emission can be interpreted as coming from the CBP footpoints \citep{2011A&A...526A.134A} while high temperature coronal emission appears compact and probably comes from loop tops. CBPs are always associated with photospheric magnetic bipolar features (MBFs) \citep[e.g.,][]{1971IAUS...43..397K, 1977SoPh...53..111G, 1993SoPh..144...15W, 2003A&A...398..775M,2016ApJ...818....9M}. \citet{1994ASPC...68..377H} analyzed simultaneous time sequence of longitudinal magnetic field data and X-ray observations, and found that two-thirds of all magnetic bipoles appear not to be related to X-ray BPs. The authors concluded that ``emergence and cancellation of magnetic flux in the photosphere is not in itself a necessary and sufficient condition for the occurrence of XBPs but rather the interaction and reconnection of magnetic field with the existing, overlying magnetic field configuration that results in the occurrence and variability of X-ray BPs''. A recent study by \citet{2016ApJ...818....9M} on the formation of MBFs associated with CBPs using Helioseismic Magnetic Imager (HMI) on board SDO data confirmed the earlier finding by, e.g., \citet{1993SoPh..144...15W} and \citet{1985AuJPh..38..875H}, that three processes are involved in the formation and evolution of CBPs: emergence, convergence and local coalescence of photospheric magnetic concentrations. This study also established that 50\%\ of 70 CBPs are formed by bipolar flux emergence (the so-called ephemeral regions). \begin{table*} \centering \caption{General information on the QS BPs associated with eruptions.} \begin{tabular}{ccccccccccc} \hline\hline Event & Original & Location & \multicolumn{2}{c}{Bipole form.} & Flux emergence & \multicolumn{3}{c}{CBP} & Delay$^a$ \\ \cline{4-5}\cline{7-9} No. & event & x, y& Pos & Neg & Date - Time & Start date - time & End date - time & Lifetime & (min) \\ & No. & (arcsec) & & & yyyy/mm/dd (UT) & yyyy/mm/dd (UT) & yyyy/mm/dd (UT) & & \\ \hline 1 & 4 & 130, -169 & A & A & 2011/01/02 16:55 & 2011/01/02 17:25 & 2011/01/03 07:37 & 14h 12m & 30 \\ \hline 2 & 5 & 60, -134 & A & A & 2011/01/01 06:30 & 2011/01/01 07:30 & 2011/01/02 00:40 & 17h 10m & 60 \\ \hline 3 & 8 & -109, -209 & A & A & 2011/01/01 19:20 & 2011/01/01 19:52 & 2011/01/02 16:40 & 20h 48m & 32 \\ \hline 4 & 10 & -89, -69 & B & B & -- & 2011/01/01 14:46 & 2011/01/02 07:07 & 16h 21m &-- \\ \hline 5 & 26 & 160, -144 & C & C & -- & 2011/01/01 01:45 & 2011/01/02 00:50 & 23h 5m & -- \\ \hline 6 & 41 & -359, 335 & A & A & 2011/01/02 10:15 & 2011/01/02 11:20 & 2011/01/03 17:00 & 29h 40m & 65 \\ \hline 7 & 50 & -379, 205 & A & A & 2010/12/31 20:00 & 2010/12/31 20:29 & 2011/01/01 20:20 & 23h 51m & 29 \\ \hline 8 & 53 & -429, 205 & A + B & B & 2011/01/01 12:50 & 2011/01/01 14:50 & 2011/01/02 01:40 & 10h 50m & -- \\ \hline 9 & 57 & -239, 310 & A & A & 2011/01/02 11:30 & 2011/01/02 12:04 & 2011/01/03 14:50 & 26h 46m & 34 \\ \hline 10 & 59 & -14, 305 & B & A + B & -- & 2011/01/03 02:30 & 2011/01/03 20:40 & 18h 10m & -- \\ \hline 11 & 61 & -254, 315 & B & A + B & 2011/01/01 22:20 & 2011/01/01 20:40 & 2011/01/02 23:30 & 26h 50m & -- \\ \hline \multicolumn{10}{l}{Notes.} \\ \multicolumn{10}{l}{A: emergence; B: convergence; C: local coalescence} \\ \multicolumn{10}{l}{$^a$: Time delay between the flux emergence and the CBP formation} \end{tabular}% \label{tab:bp}% \end{table*} CBPs at all latitudes were found to exhibit flaring activity that presents itself as an intensity increase of a few orders of magnitude in a small area of the CBPs \citep{1974ApJ...189L..93G}. \citet{1977ApJ...218..286M} associated ``flaring'' in X-ray BPs with macrospicules observed in H$\alpha$, thus for the first time establishing a link between CBP micro-flaring and eruptive activity in CBPs. Later, \citet{2009A&A...495..319I} introduced the term mini coronal mass ejections (mini-CMEs) as small-scale eruptions from the quiet Sun which appear conspicuously similar to CMEs in on-disk 171~\AA\ EUV image sequences. To recall, CMEs are large-scale eruptive phenomena. In white light observations, a classic CME is composed of three parts: a bright loop, moving ahead of a dark cavity, and a bright core which corresponds to an eruptive prominence. However, only 30\% of the CMEs are found to have this structure with some lacking a bright core due to the draining of the prominence material towards the solar surface or rather a general lack of any prominence material \citep{2011LRSP....8....1C}. The mini-CMEs were found at the junctions of supergranulation cells and the majority of the events showed cool plasma ejections that appear as darkening as referred to by the authors (e.g., mini- filament eruption) and micro-flaring in the eruption source region. Several events showed coronal dimmings that were identified as wave-like features propagating from the eruption site. Before the term mini-CME was introduced, a statistical study on EUV coronal jets by \citet{2009SoPh..259...87N} of 79 jet events in polar coronal holes identified 37 Eiffel tower-type jet events, 12 lambda-type and five micro-CME-type. The micro-CME-type jets earned their name from the great similarities with classical CMEs. The micro-CMEs appeared as small loops expanding from the solar surface. In \citet{2010A&A...517L...7I} two events, one in the quiet Sun and another in an equatorial coronal hole, were observed simultaneously on the solar disk in the extreme ultraviolet imaging telescope (EUVI)/SECCHI 171~\AA\ images of STEREO-A, and at the limb in the 304~\AA\ images from EUVI/SECCHI on board STEREO-B. The comparison of the timing of the chromospheric eruption observed at the limb in the 304~\AA\ passband with the brightening and the dimming in the disk centre in 171~\AA\ revealed that a coronal dimming appears prior to the chromospheric eruption in both events. The authors suggested that the removal of the overlying coronal field is fundamental in triggering these phenomena similar to CMEs. \citet{2010ApJ...709..369P} also analysed mini-CME events in EUVI/SECCHI 171~\AA. Their results showed that the mini-CMEs are characterised by a smaller size and a shorter lifetime compared with their large-scale counterparts. Small-scale coronal waves (CWs) and coronal dimmings are both observed in their events. The speeds of the coronal wave generated from mini-CMEs were approximately 10--20 times smaller than large-scale ones. Furthermore, they found that the small-scale coronal dimmings have two types: deep core dimming and a more widespread dimming, which is very similar to the dimming in large-scale CMEs. None of the above studies investigated mini-CME events in the context of CBPs. Small-scale filamentary (often appearing as small loops) eruptive phenomena found in H$\alpha$ images of the quiet Sun were explored by \citet{1986NASCP2442..369H} using data from the Big Bear Solar Observatory (BBSO). The features had an average length of 15\arcsec\ and were found to have an occurrence rate of 600 per 24~hrs with an average duration of the eruption of 24~min from an average lifetime of the mini-filaments of 70~hrs. The authors note that the eruptions usually proceed as an expansion of the mini-filament (MF) into an arch or a loop. All events were related to small flares. Importantly, the filaments were associated with canceling magnetic bipoles. The authors state that the MFs were either located at the bipole inversion line, or one or both footpoints of the MF were rooted there. A dedicated CBP study by \citet{2014ApJ...796...73H} used data from the Atmospheric Imaging Assembly (AIA) on board the Solar Dynamic Observatory (SDO) to study the eruptive behaviour of 30 CBPs in a coronal hole. One-quarter to one-third of all CBPs were found to have one or more mini-filament eruptions. The authors concluded that convergence and cancellation of magnetic dipoles associated with the CBPs are possibly responsible for the erupting MFs confirming the conclusions of \citet{1986NASCP2442..369H}. By exploring the state-of-the-art capabilities of the present solar observations and modelling, we aim to obtain in full detail the nature and evolutionary characteristics of eruptions from quiet Sun CBPs, by analysing AIA/SDO images in four channels, He~{\sc ii}~304~\AA, Fe~{\sc ix/x} 171~\AA, Fe~{\sc xii}~193~\AA, and Fe~{\sc xviii} 94~\AA, together with longitudinal magnetograms from HMI/SDO. The derived observational properties are used in the second part of this study where a data-driven modelling based on a relaxation code is used to model the time evolution of the analysed CBPs. The article is organised as follows: In Sect.~\ref{sect:obs} we describe the observations and the methodology. The results are presented in Sect.~\ref{sect:res}. The discussion is given in Sect.~\ref{sect:sum} and the conclusions in Sect.~\ref{sect:concl}. | \label{sect:concl} The observational evidence presented here indicates that the majority of coronal bright points (CBPs) in the quiet Sun produce mass ejections during their lifetime. The eruptions usually occur late in the evolution of the CBPs during the convergence and cancellation phase of the CBP bipole evolution during which the CBPs become smaller until they fully disappear. Eruptions can also occur earlier when CBP bipoles emerge in regions with a strong pre-existing flux. For the majority of the eruptions magnetic convergence and cancellation involve the CBP main bipoles, while in just three eruptions one of the CBP magnetic fragments and a preexisting fragment of opposite polarity converge and cancel. The CBP eruptions appear as expulsions of chromospheric material either as elongated filamentary structure (mini-filament, MF) or a volume of cool material (cool plasma cloud, CPC), together with the CBP or higher overlying hot loops. Coronal waves are occasionally identified during the eruptions. A micro-flaring is observed beneath all erupting MFs/CPCs (one case is inconclusive). It remains uncertain whether the destabilised MF causes the micro-flaring or the destabilisation and eruption of the MF is triggered by reconnection beneath the filament. In most eruptions, the cool erupting plasma obscures partially or fully the micro-flare until the erupting material moves away from the CBP. Around 50\% of the eruptions represent mini-CMEs. The dimming regions associated with the CMEs are occupied by both the `dark' cool plasma and areas of weakened coronal emission caused by the depleted plasma density. The present study demonstrates that the evolution of small-scale loop structures in the quiet Sun determined by their magnetic footpoint motions and/or ambient field topology evolve into an eruptive phase that triggers the ejection of cool and hot plasma in the corona. In our follow-up study we model the time evolution of the magnetic field, using a potential extrapolation of the first HMI magnetogram in the time series. The 3D magnetic field structure is then advanced in time assuming a relaxation approach where the photospheric boundary conditions of the magnetic field is changed according to the observed HMI magnetograms. The evolution of the magnetic fields in the CBP regions are then investigated to learn more about the general magnetic field topologies and the possibilities for building up structures that may resemble the eruption configurations. | 18 | 8 | 1808.04541 |
1808 | 1808.04294_arXiv.txt | We report the discovery with the {\it Neutron Star Interior Composition Explorer} ({\it NICER}) of mHz X-ray brightness oscillations from the ``clocked burster'' \sourcefull{}. {\it NICER} observed the source in the periods 2017 June 20 - 29, July 11 - 13, and September 9 - 15, for a total useful exposure of 34 ks. Two consecutive dwells obtained on 2017 September 9 revealed highly significant oscillations at a frequency of 8 mHz. The fractional, sinusoidal modulation amplitude increases from $0.7 \%$ at 1 keV to $\approx2\%$ at 6 keV. Similar oscillations were also detected at lower significance in three additional dwells. The oscillation frequency and amplitude are consistent with those of mHz QPOs reported in other accreting neutron star systems. A thermonuclear X-ray burst was also observed on 2017 June 22. The burst properties and X-ray colors are both consistent with \source{} being in a soft spectral state during these observations, findings that are confirmed by ongoing monitoring with {\it MAXI} and {\it SWIFT}-BAT. Assuming that the mHz oscillations are associated with black body emission from the neutron star surface, modeling of the phase-resolved spectra shows that the oscillation is consistent with being produced by modulation of the temperature component of this emission. In this interpretation, the black body normalization, proportional to the emitting surface area, is consistent with being constant through the oscillation cycle. We place the observations in the context of the current theory of marginally stable burning and briefly discuss the potential for constraining neutron star properties using mHz oscillations. | \label{sec:introduction} Neutron stars in accreting X-ray binaries are known to exhibit phenomena associated with unstable nuclear burning on their surfaces. At present we know of 111 such systems\footnote{https://personal.sron.nl/\textasciitilde jeanz/bursterlist.html} \citep{2017arXiv171206227G}. These objects produce hydrogen and helium-powered thermonuclear X-ray flashes, also known as Type I X-ray bursts \citep{2006csxs.book..113S}, as well as the longer and rarer ``superbursts,'' likely powered by carbon burning \citep{2002ApJ...566.1045S, 2001ApJ...559L.127C}. A much smaller number of objects, five at present count, have shown low-frequency ($\approx 7-10$ mHz) oscillations very likely associated with the transition from stable to unstable burning, known as marginally stable nuclear burning \citep{2007ApJ...665.1311H, 2017arXiv171206227G}. \cite{2001A&A...372..138R} reported the first detections of mHz quasiperiodic oscillations (QPOs) in three accreting neutron star systems, 4U 1608$-$52, 4U 1636$-$536, and Aql X-1. They argued that the mHz oscillations were very likely associated with the neutron star surface and not the accretion flow. They also showed that they only occur within a narrow range of X-ray luminosity. Since then, mHz oscillations have also been reported from 4U 1323$-$619 \citep{2012AAS...21924903S}, and the 11 Hz pulsar IGR J17480$-$2446 in the globular cluster Terzan 5 \citep{2012ApJ...748...82L}. We note that this object showed different behavior in its mHz oscillations compared to the others. During its outburst in 2011, as the flux increased, the recurrence time of bursts steadily decreased to about 3 minutes, finally appearing as oscillations in the light curve with this period. These oscillations also appeared at a much higher inferred accretion rate than in the other mHz sources \citep{2012ApJ...748...82L, 2017arXiv171206227G}. \cite{2002ApJ...567L..67Y} found that the frequency of kilohertz QPOs are anticorrelated with the brightness of the mHz oscillations in 4U 1608$-$52, and because this relationship is opposite to the positive correlation of kHz QPO frequency with X-ray luminosity typically seen, argued that the flux variation associated with the mHz QPO must arise on the neutron star surface and not within the accretion flow. \cite{2008ApJ...673L..35A} reported on extensive {\it Rossi X-ray Timing Explorer} ({\it RXTE}) observations of 4U 1636$-$536 that revealed a close connection between the mHz oscillations and the occurrence of X-ray bursts. They found that, during intervals of mHz oscillations, as the frequency drifted down and dropped below about 7.5 mHz, the oscillations faded and X-ray bursting resumed, and, similarly to \cite{2001A&A...372..138R}, they found no mHz oscillations immediately following the X-ray bursts. More recently, \cite{2015MNRAS.454..541L} also studied episodes of mHz QPOs in 4U 1636$-$536 using {\it XMM-Newton} and {\it RXTE} data. Similarly to \cite{2008ApJ...673L..35A}, they found frequency drift of the mHz oscillations, and suggested the drift time-scale may be set by cooling of the deeper layers, as was previously argued by \cite{2009A&A...502..871K} based on hydrodynamic calculations of helium burning with rotational mixing. \cite{2016MNRAS.463.2358L} found that the bursts that occurred immediately after an episode of mHz oscillations preferentially showed a rising light curve shape with so-called positive ``convexity.'' Such bursts show fast rise times, and this has been linked to a burst ignition location at or near the star's rotational equator \citep{2008MNRAS.383..387M, 2007ApJ...657L..29C}. Thus, \cite{2016MNRAS.463.2358L} argued that these bursts, and the burning responsible for the mHz QPOs, may occur at the neutron star equator. \cite{2016ApJ...831...34S} reported results from phase-resolved spectroscopy of the mHz oscillations in 4U 1636$-$536 obtained with {\it XMM-Newton}. They found an approximately constant color temperature with pulse phase of $\approx 0.7$ keV that they associated with the neutron star surface, and argued that the oscillations result from a variation in the surface emitting area and not the temperature. They derived a surface emitting area consistent with that expected for a neutron star. The physics relevant to marginally stable burning was recognized early on by \cite{1983ApJ...264..282P}, and has been further elucidated by \cite{2007ApJ...665.1311H}. As the helium burning stability boundary is approached, the temperature dependence of the nuclear heating rate almost precisely balances that of the cooling rate. The result is a slow, quasiperiodic mode of burning, with the oscillation period given approximately by the geometric mean of the accretion and thermal time-scales. For typical neutron star parameters and the relevant accretion rates, this gives an oscillation period close to 2 min, which is in good agreement with the range of oscillation frequencies reported for the mHz oscillations. Theory predicts that for unstable burning associated with helium ignition, the boundary between stable and unstable burning should occur at local accretion rates close to the Eddington rate \citep{1998ASIC..515..419B}. The fact that observational indications based on estimates of the X-ray luminosity suggest the onset of stability at accretion rates closer to $\approx1/10$ of the Eddington rate remains a major puzzle \citep{1988MNRAS.233..437V, 2003A&A...405.1033C}. It has been suggested that confinement of the accreted fuel, and/or rotational mixing of the fuel to greater column depths could alleviate this apparent discrepancy \citep{2009A&A...502..871K}. The fact that the mHz oscillations associated with marginal stability occur at a well defined local accretion rate makes them, in principle, a particularly important probe of the accretion rate. \cite{2014ApJ...787..101K} found that the range of accretion rates in which marginally stable oscillations can occur was also sensitive to the nuclear reaction rates. Moreover, \cite{2007ApJ...665.1311H} showed that the oscillation frequency associated with marginal stability is sensitive to the hydrogen mass fraction in the accreted fuel, as well as the surface gravity. Thus, a more detailed understanding of these oscillations could lead to new probes of these quantities. The accreting neutron star binary \sourcefull{} (hereafter, \source{}) is in many ways a prototype, having earned the moniker ``clocked burster,'' for producing bursts with extreme regularity \citep{1999ApJ...514L..27U}. The recurrence times of these regular bursts were shown to match well the predictions of theory for mixed H/He bursts \citep{2004ApJ...601..466G}, and the light curve shapes were accurately modeled using theoretical calculations of rapid-proton (rp)-process burning of approximately solar composition fuel \citep{2007ApJ...671L.141H}, hence, it has also been referred to as the ``text-book'' burster. Indeed, the bursts were so uniform in their properties that \cite{2012ApJ...749...69Z} used them as ``standard candles'' to place constraints on the mass and radius of the neutron star. The ``clocked'' bursts all occurred, however, while \source{} was in a ``hard'' spectral state. This is typically diagnosed with the use of X-ray colors \citep{2006csxs.book...39V}. Based on this, \source{} is classified as an ``atoll'' source, and the hard state is also referred to as the ``island'' state. Interestingly, \cite{2016ApJ...818..135C} recently reported on observations of \source{} in 2014 June with {\it Swift} and {\it NuSTAR} during a soft spectral state, including the detection of several X-ray bursts. They found that in this spectral state the X-ray burst recurrence times were no longer regular, and the bursts themselves also differed from the hard state bursts, being generally shorter in duration. These findings are indicative of a lower hydrogen fraction in the fuel, and suggest an additional source of stable burning occurs at the higher accretion rates in the soft state. They also detected the first photospheric radius expansion (PRE) burst from \source{}, and used it to estimate the distance as $5.7 \pm 0.2$ kpc (assuming isotropic emission). Since then, the source has mostly remained in a soft state. In this paper we report the discovery with {\it NICER} of mHz oscillations from \source{} that are very likely associated with marginally stable burning. The plan of the paper is as follows. We begin with a description of the observations and present the detection of mHz oscillations, and an X-ray burst. Next, we explore the source accretion state by studying the X-ray colors, showing that the source was in a soft, ``banana'' state in the color - color diagram \citep{1989A&A...225...79H}. Next we study the average mHz pulse profile, and present the results of phase-resolved spectroscopy. We conclude with a brief summary and discussion of the implications of our findings for models of marginally stable burning and for constraining the neutron star properties of \source{}. | We report the detection with {\it NICER} of mHz oscillations in \source{} for the first time. The oscillation properties are generally consistent with those of the mHz QPOs observed in 4U 1608$-$52, 4U 1636$-$536, Aql X-1, and 4U 1323$-$619. The {\it NICER} data were obtained while \source{} was in a soft spectral state, as indicated by both {\it MAXI} \citep{2011PASJ...63S.635S} and {\it Swift/BAT} \citep{2013ApJS..209...14K} long-term monitoring from about 2015 December to the present. Indeed, the current long-term trend appears to be a slow but steady increase in the 2 - 10 keV flux (based on the {\it MAXI} data), and very weak hard X-ray flux (from {\it Swift}-BAT). Our spectral color analysis appears consistent with this conclusion as well. The observed behavior (Figures 5, 6 and 7) appears consistent with \source{} tracing out a portion of the so-called ``banana'' branch in the color - intensity diagram during the {\it NICER} observations. The mass accretion rate is generally inferred to increase from left to right along the arc of the ``banana.'' To further test this we extracted spectra from several intervals which populate the lower left portion of the color - color diagram (see Figure 7), in order to estimate their fluxes and compare with the mean flux measured during the September epoch intervals with strong mHz oscillations (see \S 2.2 above). We find these spectra can be fit with the same model described above, and we find the unabsorbed flux ($0.6$ - $9$ keV) in the range from $5.0 - 5.2 \times 10^{-9}$ erg cm$^{-2}$ s$^{-1}$, supporting the conclusion that the X-ray flux, and therefore mass accretion rate, increases from lower left to upper right in the color - color diagram (Figure 7). Assuming the X-ray luminosity can be written as, \begin{equation} L_x = \frac{GM \dot M}{(1+z)R} = 4\pi d^2 f_x \; , \end{equation} where $M$, $R$, $(1+z) = (1 - 2GM/c^2 R)^{-1/2}$, and $f_x$ are the neutron star mass, radius, surface redshift, and X-ray flux respectively, we estimate an average mass accretion rate, $\dot M$, for the September dwells with mHz oscillations of $\dot M = 0.18 \dot M_{Edd}$, where we have assumed $M = 1.4 M_{\odot}$, $R = 10$ km, $d = 5.7$ kpc, and $\dot M_{Edd} = 1.7 \times 10^{-8} M_{\odot}$ yr$^{-1}$. This estimate also assumes that the accretion luminosity is radiated isotropically. Previous studies have shown that mHz oscillations can be associated with bursting activity. For example, \cite{2008ApJ...673L..35A} found episodes of mHz oscillations in 4U 1636$-$536 with decreasing frequency, and when the frequency dropped below about 7.5 mHz the oscillations faded and bursting resumed; however, such behavior was not associated with {\it all} X-ray bursts (see their Figure 3). Indeed, the data for 4U 1636$-$536 suggest that intervals of mHz oscillations can come and go without an obvious trigger, as long as the source is in the appropriate accretion state, indicated by its position on the banana branch or in the transition to the banana branch (see Altamirano et al. 2008, Figure 1). For \source{} we find no evidence for mHz oscillations after the X-ray burst on 22 June, which at face value is consistent with the behavior seen in 4U 1636$-$536. However, we emphasize that the present data are very sparse around the time of the X-ray burst. The only thing we can say with certainty is that we did not detect mHz oscillations in any of the {\it NICER} dwells preceeding the X-ray burst or immediately after it. Our results suggest the following long term behavior in \source{}. Prior to 2014 the source was almost exclusively in the hard state, with typical accretion rates lower than indicated by the present {\it NICER} observations. For example, comparing long term variations in the persistent X-ray flux--a reasonable proxy for mass accretion rate--we find from \cite{2004ApJ...601..466G} that the $2.5 - 25$ keV persistent (absorbed) flux from {\it RXTE} observations was $2.2 \times 10^{-9}$ erg cm$^{-2}$ s$^{-1}$ on July 29, 2002. Based on the {\it NICER} spectral modeling for the 2017 September epochs described above we estimate a flux in the same band of $4.5 \times 10^{-9}$ erg cm$^{-2}$ s$^{-1}$, a factor of two or so higher. In making this estimate we note that we have extrapolated the best-fit {\it NICER} model (Table 1) to 25 keV, that is, outside of {\it NICER}'s nominal bandpass. However, the (absorbed) flux measured by {\it NICER} in the $2.5 - 9$ keV band is $3.6 \times 10^{-9}$ erg cm$^{-2}$ s$^{-1}$, which is already significantly larger than the $2.5 - 25$ keV {\it RXTE} flux. Since the $2.5 - 9$ keV {\it NICER} flux measurement is a firm lower bound to the $2.5 - 25$ keV flux, this strongly suggests that the flux measured by {\it NICER} in 2017 September {\it in the same band} is significantly higher than the July 2002 flux measured with {\it RXTE}. This provides further support to the notion that the apparent long-term increase in X-ray flux likely represents an increasing mass accretion rate. While in the hard state the source was extensively observed with {\it RXTE}, but no mHz QPOs were ever observed. As the source continued to brighten and soften, the accretion rate became large enough to enter the accretion regime where mHz oscillations are observed in other LMXBs (see, for example, Altamirano et al. 2008, Figure 1). This overall behavior appears consistent with the observations of the mHz QPOs in 4U 1636$-$536, 4U 1608$-$52 and Aql X-1, and suggests we are seeing the same phenomenon in \source{}. More observations of \source{} in its soft, higher accretion rate state will be needed to determine if it shows the full range of mHz oscillation phenomenology evident in other sources, including a closer connection with X-ray bursts. There have been some reports of QPOs not linked with marginally stable burning, but with similar frequencies. In some of those cases the accretor is known or strongly suspected to be a black hole and not a neutron star. Examples are the QPOs identified in the black hole candidates H1743$-$322 \citep{2012ApJ...754L..23A}, and LMC X-1 \citep{2014MNRAS.445.4259A}. In the case of H1743$-$322, the QPOs are seen in a hard spectral state, whereas the mHz QPOs reported here occur in the soft state of \source{}. For LMC X-1 the observed QPO frequencies are higher, at $\approx 27$ mHz, than any previously reported mHz oscillation associated with marginally stable burning. Based on this it appears unlikey to us that the mHz oscillations observed in \source{} are related to these QPOs. Several accreting millisecond X-ray pulsars have also occasionally shown episodes of QPOs with mHz frequencies during some outbursts, an example being the 8 mHz ``flaring'' observed in IGR J00291+5934 \citep{2017MNRAS.466.3450F}. In those QPOs, however, the modulation amplitude was $> 10\%$ (rms) for photon energies less than 2 keV, which is substantially larger than is found for the mHz oscillations linked to marginally stable burning. Morever, the energy spectrum softened at the peak of these flares, which is opposite to the behavior indicated in \source{}. This argues rather strongly that the mHz oscillations reported here from \source{} are not related to these mHz ``flaring'' QPOs. As discussed in \S 2.1 the oscillations reported here from \source{} have a rather high quality factor, that is, they appear to be rather coherent over the intervals they are observed. We emphasize, however, that these intervals ($\approx 2$ ks) are relatively short compared to the hours-long time-scales over which mHz QPO frequencies have been observed to drift in several sources \citep{2015MNRAS.454..541L, 2008ApJ...673L..35A}. Additionally, current modeling of marginally stable burning can produce quite regular trains of mHz pulses, with the limiting factor here likely being temporal variations of the mass accretion rate \citep{2007ApJ...665.1311H, 2014ApJ...787..101K}. Based on these arguments we think that the relatively high coherence seen in the mHz oscillations reported here is not inconsistent with their connection to marginally stable burning. Figure 9 shows the energy dependence of the fractional (sinusoidal) modulation amplitude of the mHz oscillations in \source{}. From about $1 - 3$ keV an increasing trend is rather clear, with an apparent flattening above 3 keV. This appears to be at least qualitatively similar to the behavior seen in IGR J17480$-$2446--which shows an increasing trend in rms amplitude from about $2.5 - 10$ keV based on {\it RXTE} data \citep{2012ApJ...748...82L}--but opposite to that reported for 4U 1608$-$52 and 4U 1636$-$536 by \cite{2001A&A...372..138R}, also using {\it RXTE}. However, some caution is warranted here due to the different bandpasses of {\it NICER} and {\it RXTE}, as well as the extent to which uncertainties in backgrounds could influence derived amplitudes. A closer comparison of the energy dependence of the mHz oscillations reported here with {\it NICER}, and those detected with {\it XMM-Newton} should prove useful, as these instruments have more closely matched bandpasses. We extracted spectra as a function of mHz pulsation phase, and find that they are consistent with the modulation being produced by a variation in the temperature of a thermal (black body) component at an effectively constant emitting area, though we emphasize again that this is not a unique spectral interpretation (see discussion in \S 3 above). Assuming a spherical emitting source and isotropic emission, the fitted black body normalization implies a radius of $R \approx 14$ km. Given uncertainties in the source distance, the exact nature of the continuum describing the persistent emission, the extent to which disk reprocessing is present, as well as our neglect of color corrections, this value should only be taken as roughly indicative of the surface area associated with the mHz oscillations, and definitely not a precise measurement. The evidence that the oscillating flux can be associated with a thermal (black body) component with an area approximately consistent with that for a neutron star provides support for the interpretation of the oscillations as due to marginally stable nuclear burning; however, more data will be needed to establish this definitively. We note that \cite{2016ApJ...831...34S} reached different conclusions from a study of the phase-resolved spectra of mHz oscillations from 4U 1636$-$536. They argued that the flux modulation resulted from a variation in the emitting area at approximately constant temperature. While this is at odds with current theory, which finds that the temperature oscillates \citep{2007ApJ...665.1311H}, we emphasize that current theoretical calculations of marginally stable burning are all one dimensional, so they essentially do not explore the possibility of variations in the burning area. Here we point out that there are several physical effects that suggest that lateral variations in the burning should be included in future modeling. For example, \cite{2007ApJ...665.1311H} showed that the oscillations associated with marginal stability occur in a very narrow range of mass accretion rates, $\dot m$, for a given surface gravity. They found the oscillations were present in a range of only $1\%$ of the critical accretion rate. Note that it is the {\it local} accretion rate that is relevant for the nuclear burning. For a fast spinning neutron star the variations in effective surface gravity from equator to pole can result in a latitudinal change in the local accretion rate. \cite{2007ApJ...657L..29C} used this effect to explore the extent to which ignition of thermonuclear bursts could preferentially occur at different latitudes. However, this process would appear to be relevant for marginally stable burning as well, as a greater than $1\%$ variation in the local accretion rate with latitude could influence the range of stellar latitudes for which marginally stable oscillations could occur. Indeed, this appears to be a possible mechanism for limiting the marginally stable burning to a belt in latitude \citep{2016MNRAS.463.2358L}. As a rough guide, a star rotating at 582 Hz (close to the inferred spin frequency for 4U 1636$-$536), and with a mass of $1.4 M_{\odot}$ and equatorial radius of 10 km, would have an $\approx 11\%$ increase in effective surface gravity from equator to pole, using the results of \cite{2014ApJ...791...78A}. This is certainly large enough to have a significant effect on the local accretion rate at the level required to influence marginally stable burning. Interestingly, for a spin rate of 11 Hz, appropriate for IGR J17480$-$2446 in Terzan 5, the variations in surface gravity are $< 0.004 \%$ equator to pole, and thus latitudinal variations would likely be negligible for such a slow spinner. This could perhaps account for some of the observed differences in the mHz oscillations of IGR J17480$-$2446 and fast spinners like 4U 1608$-$52, 4U 1636$-$536, and Aql X-1. As pointed out by \cite{2007ApJ...665.1311H}, the critical accretion rate for marginally stable oscillations is itself a function of surface gravity, so the occurrence of oscillations, and their properties, will in principle depend on both effects; however, this discussion highlights that such latitudinal dependencies of the burning are likely, at least for fast spinners, and should be considered in future modeling of marginally stable oscillations. \cite{2007ApJ...665.1311H} argued that the oscillation frequency associated with marginally stable burning should be sensitive to the neutron star surface gravity, which is proportional to $M/R^2$, and the mass fraction of hydrogen in the accreted matter (see their Figure 9). While their results on this score were based on one-zone models, the good overall agreement between predictions of their one-zone models and multi-zone hydrodynamic calculations with the KEPLER code \citep{2007ApJ...665.1311H}, suggests that these dependencies are robust. This is particularly promising for \source{} because of the extensive theoretical modeling already done on the ``clocked'' bursts \citep{2007ApJ...671L.141H, 2018ApJ...860..147M}, which suggests a near solar composition for the accreted fuel. Interestingly, the bursts observed so far at the higher accretion rates associated with the soft spectral state have shorter durations, and several have reached higher peak fluxes, as is evidenced by the occurrence of PRE. These are both good indications of depletion of hydrogen in the fuel compared to the hard-state, ``clocked'' bursts. This suggests that additional theoretical modeling of the soft state bursts, coupled with further comparisons to the ``clocked'' bursts, could place tighter constraints on the hydrogen fraction present in the fuel during episodes of mHz oscillations. It might then be possible to provide more robust limits on the neutron star surface gravity from measurements of the mHz oscillations. | 18 | 8 | 1808.04294 |
1808 | 1808.06128_arXiv.txt | We present a single-dish mapping algorithm with a number of advantages over traditional techniques. (1)~Our algorithm makes use of weighted modeling, instead of weighted averaging, to interpolate between signal measurements. This smooths the data, but without blurring the data beyond instrumental resolution. Techniques that rely on weighted averaging blur point sources sometimes as much as 40\%. (2)~Our algorithm makes use of local, instead of global, modeling to separate astronomical signal from instrumental and/or environmental signal drift along the telescope's scans. Other techniques, such as basket weaving, model this drift with simple functional forms (linear, quadratic, etc.\@) across the entirety of scans, limiting their ability to remove such contaminants. (3)~Our algorithm makes use of a similar, local modeling technique to separate astronomical signal from radio-frequency interference (RFI), even if only continuum data are available. (4)~Unlike other techniques, our algorithm does not require data to be collected on a rectangular grid or regridded before processing. (5)~Data from any number of observations, overlapping or not, may be appended and processed together. (6)~Any pixel density may be selected for the final image. We present our algorithm, and evaluate it using both simulated and real data. We are integrating it into the image-processing library of the Skynet Robotic Telescope Network, which includes optical telescopes spanning four continents, and now also Green Bank Observatory's 20-meter diameter radio telescope in West Virginia. Skynet serves hundreds of professional users, and additionally tens of thousands of students, of all ages. Default data products are generated on the fly, but will soon be customizable after the fact. | \subsection{Skynet} Founded in 2005, Skynet is a global network of fully automated, or robotic, volunteer telescopes, scheduled through a common web interface.\footnote{https://skynet.unc.edu} Currently, our optical telescopes range in size from 14 to 40 inches, and span four continents. Originally envisioned for gamma-ray burst follow-up (Reichart et al.\@ 2005, Haislip et al.\@ 2006, Dai et al.\@ 2007, Updike et al.\@ 2008, Nysewander et al.\@ 2009, Cenko et al.\@ 2011, Cano et al.\@ 2011, Bufano et al.\@ 2012, Jin et al.\@ 2013, Morgan et al.\@ 2014, Martin-Carrillo et al.\@ 2014, Friis et al.\@ 2015, De Pasquale et al.\@ 2016, Bardho et al.\@ 2016, Melandri et al.\@ 2017), Skynet has also been used to study gravitational-wave sources (Abbott et al.\@ 2017a, 2017b, Valenti et al.\@ 2017, Yang et al.\@ 2017), blazars (Osterman Meyer et al.\@ 2008, Valtonen et al.\@ 2016, Zola et al.\@ 2016, Liu et al.\@ 2017, Goyal et al.\@ 2017), supernovae (Foley et al.\@ 2010, 2012, 2013, Pignata et al.\@ 2011, Valenti et al.\@ 2011, 2014, Pastorello et al.\@ 2013, Milisavljevic et al.\@ 2013, Maund et al.\@ 2013, Fraser et al.\@ 2013, Stritzinger et al.\@ 2014, Inserra et al.\@ 2014, Takats et al.\@ 2014, 2015, 2016, Dall'Ora et al.\@ 2014, Folatelli et al.\@ 2014, Barbarino et al.\@ 2015, de Jaeger et al.\@ 2016, Gutierrez et al.\@ 2016, {\color{black}2018,} Tartaglia et al.\@ 2017, 2018, Prentice et al.\@ {\color{black}2018}), supernova remnants (Trotter et al.\@ 2017), novae (Schaefer et al.\@ 2011), pulsating white dwarfs and hot subdwarfs (Thompson et al.\@ 2010, Barlow et al.\@ 2010, 2011, 2013, 2017, Reed et al.\@ 2012, Bourdreaux et al.\@ 2017, Hutchens et al.\@ 2017), a wide variety of variable stars (Layden et al.\@ 2010, Gvaramadze et al.\@ 2012, Wehrung et al.\@ 2013, Miroshnichenko et al.\@ 2014, Abbas et al.\@ 2015, Khokhlov et al.\@ 2017, 2018), a wide variety of binary stars (Reed et al.\@ 2010, Sarty et al.\@ 2011, Helminiak et al.\@ 2011, 2012, 2015, Strader et al.\@ 2015, Tovmassian et al.\@ 2016, 2017, Fuchs et al.\@ 2016, Pala et al.\@ 2017, Neustroev et al.\@ 2017, Lubin et al.\@ 2017, Swihart et al.\@ 2017, Zola et al.\@ 2017, Kriwattanawong et al.\@ 2018), exoplanetary systems (Fischer et al.\@ 2006, Czesla et al.\@ 2012, Meng et al.\@ 2014, 2015, Kenworthy et al.\@ 2015, Kuhn et al.\@ 2016, Awiphan et al.\@ 2016, Blank et al.\@ 2018), trans-Neptunian objects and Centaurs (Braga-Ribas et al.\@ 2013, 2014, Dias-Oliveira et al.\@ 2015), asteroids (Descamps et al.\@ 2009, Pravec et al.\@ 2010, 2012, 2016, Marchis et al.\@ 2012, Savanevych et al.\@ 2018), and near-Earth objects (NEOs; Brozovic et al.\@ 2011, Pravec et al.\@ 2014). Skynet is also the leading tracker of NEOs in the southern hemisphere (R. Holmes, private communication). \begin{figure} \plotone{fig1.pdf} \caption{Green Bank Observatory 20-meter diameter radio telescope. (Photo credit: GBO)} \end{figure} \begin{figure*} \plotone{fig2.pdf} \caption{OTF mapping patterns. \textbf{Left:} Raster (horizontal). \textbf{Middle:} Nodding. \textbf{Right:} Daisy.} \end{figure*} Skynet's mission is split evenly between supporting professional astronomers and supporting students and the public. Although most of our observations have been for professionals, most of our users are students. We have developed/are continuing to develop Skynet-based curricula for undergraduates, high-school students (in partnership with Morehead Planetarium and Science Center (MPSC)), and middle school-aged students (in partnership with the University of Chicago/Yerkes Observatory, Green Bank Observatory (GBO), the Astronomical Society of the Pacific, and 4-H), as well as for blind and visually-impaired students (in partnership with Associated Universities, Inc., the University of Chicago/Yerkes Observatory, the Technical Education Research Centers, and the University of Nevada at Las Vegas). These efforts have reached over 20,000 students, and our public-engagement efforts (also in partnership with MPSC) have also reached over 20,000, mostly elementary and middle school-aged students. Curriculum-based student users queue observations through the same web interface that the professionals use. Altogether, over 15 million images have been taken to date. In partnership with GBO, and funded by the American Recovery and Reinvestment Act, Skynet has added its first radio telescope, GBO's 20-meter in West Virginia (see Figure~1). We describe the 20-meter, which has been refurbished, in \textsection2. As with Skynet's optical telescopes, the 20-meter serves both professionals and students. Professional use consists primarily of timing {\color{black}observations (e.g., pulsar timing; fast radio burst searches in conjunction with \textit{Swift})}, but also some mapping observations{\color{black}, for photometry} (e.g., {\color{black}fading of Cas~A and improved flux-density calibration of the radio sky (}Trotter et al.\@ 2017{\color{black}); intra-day variable blazar campaigns in conjunction with other radio, and optical, telescopes}). Student use consists of {\color{black}timing, spectroscopic (e.g., Williamson et al.\@ 2018), and mapping observations}, but with an emphasis on mapping, at least for beginners. Regarding student, as well as public, use, the 20-meter represents a significant opportunity for radio astronomy. Small optical telescopes can be found on many, if not most, college campuses. But small radio telescopes are significantly more expensive to build, operate, and maintain, and consequently are generally found only in the remote locations that make the most sense for professional use. Consequently, most people -- including most students of astronomy -- never experience radio telescopes, let alone use them. However, under the control of Skynet, the 20-meter is not only more accessible to more professionals, it is already being used by thousands of students per year, of all ages, as well as by the public. \subsection{Single-Dish Mapping} In this paper, we present the single-dish mapping algorithm that we developed for Skynet, both for professional use and for student use. Here we outline the design requirements that we set for ourselves, and outline approaches that we adopted/developed to meet these requirements. \subsubsection{Mapping Pattern} Many single-dish mapping algorithms (e.g., Sofue \& Reich 1979, Emerson \& Grave 1988) require the signal to be sampled on a rectangular grid. Generally, this requires what is called ``step-and-integrate'' or ``point-and-shoot'' mapping. However, this is an inefficient way to observe, with greater telescope acceleration, deceleration, and settling overheads, greater wear and tear on the telescope, and greater sensitivity to time-variable systematics (see \textsection1.2.3). ``On-the-fly'' (OTF) mapping, in which signal is integrated while the telescope moves, minimizes these concerns, as long as integrations span no more than $\approx$0.2 beamwidths (FWHM) along the telescope's direction of motion (along the telescope's “scan”), to avoid blurring point sources by more than $\approx$1\% (Mangum, Emerson \& Greison 2007). A wide variety of OTF mapping patterns might be employed (e.g., see Figure~2): \begin{itemize} \item A ``raster'' pattern approximates a rectangular grid, though integrations might not line up from scan to scan. \item A ``nodding'' pattern, in which the telescope moves up and down, typically in elevation, as Earth's rotation carries the telescope's beam across the sky in the perpendicular direction, is typical of meridian-transit telescopes, which cannot move east-west (e.g., see \textsection2.2). However, even with full-motion telescopes, a nodding pattern can be used to maintain a constant parallactic angle, if it is desired that the telescope's beam pattern not rotate across the image, or from image to image. \item ``Spiral'' and ``hypocycloid'' patterns are efficient ways of mapping sources and extended structures, respectively, without abrupt changes to the telescope's motion (Mangum, Emerson \& Greison 2007). \item A ``daisy'' pattern can also be used to map sources, with the advantage of crossing the source's peak, and sampling the background level, an arbitrarily large number of times. \end{itemize} As long as the gaps between scans do not exceed Nyquist sampling, or $\approx$0.4 beamwidths, in theory all information between scans can be recovered (modulo noise, interference, etc.\@) Given this, and not wishing to limit our users to a particular mapping pattern, or set of patterns, we require that our algorithm work independently of mapping pattern. Furthermore, since student/quick-look maps may be undersampled, we also require that our algorithm work independently of sampling density (within reason). Instead, undersampled images will be flagged as not for professional use. \subsubsection{Signal Averaging vs.\@ Signal Modeling} Some single-dish mapping algorithms get around the problem of not collecting data on a rectangular grid by resampling the data onto a rectangular grid before processing (e.g., Winkel, Floer \& Kraus 2012). This is known as ``regridding'', and is typically done by taking a weighted average of the data around each grid point. This is essentially a convolution of the data with the weighting function, sometimes called the convolution kernel. However, this blurs the image. Many adopt a kernel of width equal to the telescope's beamwidth, but this results in $\approx$40\% blurring and $\approx$40\% errors (near the center of the beam pattern) in the reconstructed image (see Figure~3). Others oversample so a narrower kernel can be used. For example, Winkel, Floer \& Kraus (2012) collect $\approx$3 times as much data as required by Nyquist and use a 1/2-beamwidth kernel, but this still results in $\approx$12\% blurring and $\approx$18\% errors (near the center of the beam pattern) in the reconstructed image (Figure~3). Given that single-dish maps already suffer from poor resolution (compared to interferometric maps), we do not want to degrade resolution further, simply due to processing. To this end, instead of weighted \textit{averaging}, we use weighted \textit{modeling}: Instead of averaging the data local to each pixel in the final image, using a weighting function, we fit a model to the data local to each pixel in the final image, using a weighting function (see Figure~4). As long as (1)~the model is sufficiently flexible over the scale of the weighting function, and (2)~sampling is sufficiently dense to constrain (actually, overconstrain) this model, the signal should be recoverable at any location, without blurring. For example, the weighted modeling approach that we present in \textsection3.7 is able to recover the simulated data in Figure~3 with $<$1\% errors (near the center of the beam pattern). Another requirement that we place on our algorithm is that this -- replacing the data with a locally modeled version of itself -- be the last step, not the first step. Other algorithms regrid the data before processing (e.g., contaminant removal, see \textsection1.2.3), because their processing operations work only on gridded data. However, this is a poor modeling practice: Even if weighted averaging could be done accurately, and even if weighted modeling can be done accurately, every operation on data propagates uncertainty in ways that are both difficult to understand and even more difficult to properly account for before/in the next operation. It is always preferable to operate on real data, at least for as long as possible, than it is to operate on successive approximations thereof. \subsubsection{Contaminant Removal} Single-dish mapping algorithms must address signal contaminants, which we separate into three broad categories: (1)~en-route signal drift, resulting in what is sometimes called ``the scanning effect'', (2)~radio-frequency interference (RFI), and (3)~elevation-dependent signal. Each of these is demonstrated in Figure~5. \textbf{En-Route Drift:} Even with very stable, modern receivers, the detected signal can still drift in time, due to instrumental reasons, such as $1/f$, or pink, noise, and/or due to environmental reasons, such as changing atmospheric emission, or changing ground emission spilling over the edge of the dish, particular as the dish moves. This results in low-level variations in the signal along the telescope's scans, and is most noticeable from scan to scan (the scanning effect). Components of these variations can be made to vary over shorter or longer angular scales by moving the telescope slower or faster, but this does not eliminate them. Sofue \& Reich (1979) used unsharp masking to separate en-route drift (and small-scale structure in the image) from larger-scale structure, and then modeled the en-route drift along each scan with a second-order polynomial, using sigma-clipping to avoid small-scale structure contamination. Drawbacks of this approach are: (1)~Unsharp masking uses a blurred version of the data to correct the data, resulting in a blurred image, at least in the lower signal-to-noise parts of the image (something we wish to avoid; \textsection1.2.2); (2)~Low-order polynomials may adequately model en-route drift over small angular scales, but are too simple/increasingly inaccurate over larger angular scales, such as the length of a scan; and (3)~Sigma-clipping is too crude of an outlier rejection criterion (see \textsection1.3). \begin{figure*} \epsscale{0.95} \plotone{fig3.pdf} \caption{\textbf{First Row:} Simulation of a point source sampled with a Gaussian beam pattern on a 1/5-beamwidth grid, recovered using, from left to right: 1-beamwidth weighted averaging, 2/3-beamwidth weighted averaging, 1/2-beamwidth weighted averaging, and weighted modeling, as described in \textsection3.7. \textbf{Second Row:} Residual error associated with each of these techniques. \textbf{Third Row:} Cassiopeia~A observed with one of the 20-meter's L-band linear polarization channels using a 1/5-beamwidth raster, and recovered using the above techniques. Weighted averaging fails to recover the telescope's beam pattern, which is structured. \textbf{Fourth Row:} Difference between each of these techniques and weighted modeling. Square-root and squared scalings are used in the third and fourth rows, respectively, to emphasize fainter beam structure (units are dimensionless, with one corresponding to the noise diode; see \S3.1.)} \end{figure*} \begin{figure*} \plottwo{fig4a.pdf}{fig4b.pdf} \caption{Weighted modeling of 20-meter data from the third row of Figure~3, at two representative points. Modeled surfaces span two beamwidths, but are most strongly weighted to fit the data over only the central, typically, 1/3 -- 2/3 beamwidths, as described in \textsection3.7. Only the central point (red) is retained.} \end{figure*} \begin{figure} \plotone{fig5.pdf} \caption{Raw map of Virgo~A (top), 3C~270 (center right), and 3C~273 (bottom), acquired with the 20-meter in L band, using a 1/10-beamwidth horizontal raster. Left and right linear polarization channels have been summed, partially symmetrizing the beam pattern (Figure~3). Locally modeled surface (\textsection1.2.1, see \textsection3.7) has been applied for visualization only. All three signal contaminants are demonstrated: (1)~en-route drift, the low-level variations along the horizontal scans, (2)~RFI, both long-duration, during the scan that passes through 3C~273, and short-duration, near Virgo~A, and (3)~elevation-dependent signal, toward the upper right, which was only $\approx$11$\degr$ above the horizon.} \end{figure} \begin{figure} \plotone{fig6.pdf} \caption{Green Bank Observatory 40-foot diameter radio telescope. (Photo credit: GBO)} \end{figure} Emerson \& Grave (1988) Fourier transform the data, collapsing en-route drift to a single band of near-zero spatial frequencies, perpendicular to the scan direction. They then mask it in Fourier space and transform back. Ideally, two orthogonal mappings are combined (in Fourier space), so real spatial frequencies that are masked in one image can be recovered from the other. The primary advantage of this approach is that it does not assume that en-route drift can be well-modeled by a low-order polynomial over the length of each scan. However, drawbacks are: (1)~We will not always have two orthogonal mappings; (2)~Even if we did, OTF rasters do not populate rectangular grids, requiring the data to be regridded before processing (however, \textsection1.2.2); and (3)~This technique would not work well, or at all, with other mapping patterns (e.g., noddings, daisies, etc.; \textsection1.2.1). Haslam, Quigley \& Salter (1970), Haslam et al.\@ (1974), Seiber, Haslam \& Salter (1979), and Haslam et al.\@ (1981) introduced a technique called basket-weaving, involving two mappings with intersecting scans (not necessarily orthogonal). Signal differences at the intersections are minimized, but again assuming that en-route drift can be well-modeled by a low-order polynomial over the length of each scan. The procedure also requires iteration. Winkel, Floer \& Krauss (2012) introduced an updated version that does not require iteration, but: (1)~It does require regridding (\textsection1.2.2); (2)~It works only with two, near-orthogonal mappings, so not with single mappings, and not with noddings, daisies, etc.\@ (\textsection1.2.1); and (3)~Although it does permit an arbitrarily high-order parameterization of the en-route drift over the length of each scan, in practice it becomes less effective as the number of parameters times the number of scans approaches the number of regridded pixels (which is typically significantly less than the original number of signal measurements). This limited them to second-order polynomials in the examples that they presented. \textbf{RFI:} RFI is typically localized to specific frequencies. If spectral information is available, these frequencies can be identified and masked, or the excess signal at these frequencies can be measured and subtracted off (see \textsection3.1 of Paper II, Dutton et al.\@ 2018). However, sometimes only continuum data are available (e.g., see \textsection2.2), or spectral data are available but the RFI has a continuous spectrum (e.g., lightning). In these cases, RFI can be identified only from its temporal signature. If prolonged, it might be indistinguishable from en-route drift. If localized in time, it is unlikely to occur at the same position on adjoining scans, appearing as a source that is narrower than the telescope's beamwidth, at least across scans. In most of the above references, RFI was identified by eye and excised by hand. However, given the volume, and diversity, of users that we have with the 20-meter, we need to be able to do this automatically. \textbf{Elevation-Dependent Signal:} Atmospheric emission increases as elevation decreases. And if the dish is over-illuminated, terrestrial spillover increases as elevation increases. If interested only in small-scale structures, these backgrounds can be removed with en-route drift (see \textsection3.3). However, if large-scale structures are to be retained (or added back in), these backgrounds need to be removed separately (see \textsection2.3 of Paper II, Dutton et al.\@ 2018). \begin{figure*} \plotone{fig7.pdf} \caption{Flowchart of our algorithm for contaminant-cleaning and mapping small-scale structures. Blue are the component algorithms. Green are the inputs to and outputs of these component algorithms, consisting of data, corresponding noise models, and ultimately maps. Red are user-selected scales, for separating wanted and unwanted structures, and for modeling the final surface.} \end{figure*} \begin{figure} \plotone{fig8.pdf} \caption{40-foot gain calibration data, with the noise diode first on and then off, and best-fit model. Circled points have been robust-Chauvenet rejected, including data taken during the transitions from off to on and on to off, and RFI-contaminated data. The background level increased during the calibration, but our model accounts for this: Simply averaging each level, instead of modeling each level with a line, would have underestimated the result.} \end{figure} \begin{figure} \plotone{fig9.pdf} \caption{Point-to-point noise measurement technique. \textbf{Top:} Applied to gain-calibrated 20-meter data. The circled point has been robust-Chauvenet rejected. \textbf{Bottom:} Deviations. Mean and standard deviations are measured from the non-rejected points, for each scan.} \end{figure} \begin{figure} \plotone{fig10.pdf} \caption{Corrected 1D noise measurements vs.\@ scan number for a 20-meter observation, and best-fit model. Only two points met the robust Chauvenet rejection criterion, and then only barely, which is not unusual given that intra-scan outliers have already been rejected (Figure~9). The 1D noise level increased by $\approx$20\% over the course of this observation.} \end{figure} \begin{deluxetable}{ccc} \tablewidth{0pt} \tablecaption{Minimum Recommended 1D Background-Subtraction Scale for the Telescopes and Receivers of \textsection2, in Theoretical Beamwidths} \tablehead{ \colhead{Telescope} & \colhead{Receiver} & \colhead{Scale}} \startdata 20-meter & L (HI $+$ OH)\tablenotemark{a} & 7\tablenotemark{c} \\ 20-meter & L (HI)\tablenotemark{b} & 6\tablenotemark{c} \\ 20-meter & L (OH)\tablenotemark{b} & 6\tablenotemark{c} \\ 20-meter & X & 6\tablenotemark{c} \\ 40-foot & L (HI) & 3 \enddata \tablenotetext{a}{Before August 1, 2014} \tablenotetext{b}{After August 1, 2014} \tablenotetext{c}{The 20-meter's beam pattern has a low-level, broad component, in both L and X bands, and consequently, we recommend larger background-subtraction scales here. This component was significant in L band prior to 8/1/14, as can be seen in the third row of Figure~3, as well as in Figure~5. Post 8/1/14, it was significantly reduced, but not altogether eliminated. This component corresponds to approximately 2\% -- 3\% and 4\% -- 5\% of the integrated beam pattern in L and X band, respectively. If this is not a concern, these minimum recommended 1D background-subtraction scales can be lowered to 3 and 4 theoretical beamwidths, respectively.} \end{deluxetable} \begin{figure} \plotone{fig11.pdf} \caption{\textbf{Top:} A forward-directed local background model, anchored to an arbitrary point from Figure~5, near 3C~270. Circled points are within one background-subtraction scale length, but above the model. \textbf{Middle:} Forward- and backward-directed local background models, anchored to every point in the scan. \textbf{Bottom:} Global background model, constructed from the local background models.} \end{figure} \begin{figure} \plotone{fig12.pdf} \caption{\textbf{Top:} A preliminary local background model, anchored to the same point as in Figure~11. Circled points have been iteratively rejected as too high, given the modeled noise level (\textsection3.2); the larger points were not rejected. \textbf{Middle:} Final local background models, originally, but no longer, anchored to every point in the scan. \textbf{Bottom:} Global background model, constructed from the final local background models (solid curve), and from quadratic, instead of linear, local background models (dashed curve).} \end{figure} We approach all three types of signal contaminants in the same way. First, we model them locally, not globally: We use simple parameterizations, such as first- and second-order polynomials, but we do not expect, nor need, them to hold over large angular scales. When fitting these models, we use robust, and appropriate, outlier rejection (see \textsection1.3) to separate contaminants from astronomical signal. Finally, we combine our locally-fitted models into a global model, checking the procedure against simulated data. \subsection{Robust Chauvenet Outlier Rejection} Our single-dish mapping algorithm is built upon a new outlier rejection technique, called ``robust Chauvenet rejection'', which we developed for the Skynet Robotic Telescope Network's image-processing library in general, and for this application in particular. Sigma clipping (\textsection1.2.3) is one of the simplest outlier rejection techniques, but also one of the crudest. It suffers from a number of problems, foremost of which is how to set the threshold. For example, if working with $\approx$100 data points, 2-sigma variations are expected but 4-sigma variations are not. However, if working with $\approx$10$^4$ data points, 3-sigma variations are expected but 5-sigma variations are not. Chauvenet rejection is simply sigma clipping plus a reasonable rule for setting the threshold: \begin{equation} NP(\rm{>}|z|) = 0.5, \end{equation} \noindent where $N$ is the total number of data points and $P(\rm{>}|z|)$ is the cumulative probability of being more than $z$ standard deviations from the mean, assuming a Gaussian distribution (Chauvenet 1863). However, sigma clipping, and by extension Chauvenet rejection, also suffers from the following problem: If the mean and the standard deviation are not known a priori, which is almost always the case, they must be measured from the data, and both of these quantities are sensitive to the very outliers that they are being used to reject. This limits the applicability of Chauvenet rejection to very small sample-contamination fractions. Consequently, we developed \textit{robust} Chauvenet rejection, which makes use of the mean and the standard deviation, but also makes use of increasingly robust (but decreasingly precise) alternatives, namely the median and the (half-sample; e.g., Bickel \& Fruhwith 2005) mode, and the 68.3-percentile deviation, measured in three increasingly robust ways. These quantities approximate the mean and the standard deviation, respectively, and equal them in the case of a Gaussian distribution. But they are significantly less sensitive to outliers, and are particularly effective, meaning both robust and precise, if applied in proper combination. The applicability of this technique is very broad, spanning not only astronomy and science, but all quantitative disciplines. We have submitted this technique as a companion paper (Maples et al.\@ {\color{black}2018}), and make extensive use of it in this paper, with implementation details offered in the footnotes. As such, this paper additionally serves the purpose of ``field testing'' this new technique. \begin{figure} \plotone{fig13.pdf} \caption{20-meter 1/10-beamwidth horizontal raster replaced with Gaussian random noise, of mean zero and standard deviation one. Locally modeled surface (\textsection1.2.1, see \textsection3.7) has been applied for visualization only.} \end{figure} \begin{figure*} \plotone{fig14.pdf} \caption{\textbf{Top Row:} Data from Figure~13 background-subtracted, with 6- (left), 12- (middle), and 24- (right) beamwidth scales (the map is 24 beamwidths across). \textbf{Bottom Row:} Data from the top row minus the data from Figure~13 (residuals). Background-subtracted data are not biased high nor low. To first order, the noise level of the background-subtracted data is $\approx$98.0\% (left), $\approx$98.8\% (middle), and $\approx$99.3\% (right) of that of the original data, and the RMS of the residuals is only $\approx$20.1\% (left), $\approx$15.4\% (middle), and $\approx$12.3\% (right) of the noise level of the original data (see Figure~15 for second-order effect). Locally modeled surfaces (\textsection1.2.1, see \textsection3.7) have been applied for visualization only.} \end{figure*} \begin{figure} \plotone{fig15.pdf} \caption{Gaussian random noise background-subtracted (top) and residuals (bottom) vs.\@ $W$, the sum of the weights of the non-rejected local background models that determine the global background model at each point, for 1/10- and 1/5-beamwidth rasters, and for 1-, 3-, 6-, 12-, and 24-beamwidth background-subtraction scales (background-subtraction scale times sampling density of the data sets increases from left to right). The RMS of the data varies with $W$, but not with background-subtraction scale or sampling density independently. Curves are 1-, 2-, and 3-$\sigma$ model noise envelopes that have been fitted to all of these data simultaneously (Equations 6 and 7).} \end{figure} \begin{figure} \plotone{fig16.pdf} \caption{Simulated data from Figure~13 to which we have added point sources and short-duration RFI. For the point sources, we use a Gaussian beam pattern. For the short-duration RFI, we use the absolute value of a sum of rapidly varying sine functions multiplied by a short-duration Gaussian envelope function. Locally modeled surface (§\textsection1.2.1, see \textsection3.7) has been applied for visualization only. Square-root scaling is used to emphasize fainter structures.} \end{figure} \begin{figure*} \epsscale{0.95} \plotone{fig17.pdf} \caption{\textbf{Top Row:} Data from Figure~16 background-subtracted, with 6- (left), 12- (middle), and 24- (right) beamwidth scales (the map is 24 beamwidths across). \textbf{Middle Row:} Data from the top row (1)~minus the data from Figure~16 (residuals) and (2)~minus the Gaussian random noise residuals from the bottom row of Figure~14 (for greater clarity). Small-scale structure residuals are biased negative, but typically by at most $\approx$1/2 -- 1 (left), $\approx$1/4 -- 1/2 (middle), and $\approx$1/8 -- 1/4 (right) of the noise level, and independently of the brightness of the proximal small-scale structure (point source or short-duration RFI). Larger values are possible when small-scale structures blend together into large-scale structures, in the scan direction, where the division between small and large scales is given by the background-subtraction scale. Larger values are also possible when small-scale structures occur near the ends of scans. Noise-level biases can be ignored for all but the lowest-S/N sources (see \textsection4), and are further mitigated by our RFI-subtraction algorithm in \textsection3.6.2, and by our large-scale structure algorithm in Paper II. \textbf{Bottom Row:} Same as the middle row, but for more-realistic, less-winged sources (given by Equation~9 with $\theta_{RFI} = 1$ beamwidth and $z_0 = 0$; see Figure~31); residuals are $\approx$2 -- 3 times smaller in this case. Locally modeled surfaces (\textsection1.2.1, see \textsection3.7) have been applied for visualization only. Square-root and hyperbolic-arcsine scalings are used in the top and bottom two rows, respectively, to emphasize fainter structures.} \end{figure*} \begin{figure} \plotone{fig18.pdf} \caption{Simulated data from Figure~16 to which we have added en-route drift and long-duration RFI. For the en-route drift, we use a sum of randomly phased sine functions, the shortest period of which is 12 beamwidths. We linearly increase the maximum amplitude of the en-route drift from zero times the noise level at the bottom of the image to 12 times the noise level at the top of the image. For the long-duration RFI, we use a similarly constructed sum of sine functions, but plus a constant to ensure that it is always positive, and multiplied by a long-duration Gaussian envelope. The long-duration RFI is significantly brighter than the en-route drift. Locally modeled surface (\textsection1.2.1, see \textsection3.7) has been applied for visualization only. Square-root scaling is used to emphasize fainter structures.} \end{figure} \begin{figure*} \plotone{fig19.pdf} \caption{\textbf{Top Row:} Data from Figure~18 background-subtracted, with 6- (left), 12- (middle), and 24- (right) beamwidth scales (the map is 24 beamwidths across). \textbf{Bottom Row:} Data from the top row (1)~minus the data from Figure~16 (residuals) and (2)~minus the Gaussian random noise residuals from the bottom row of Figure~14, and the small-scale structure residuals from the middle row of Figure~17 (for greater clarity). En-route drift and long-duration RFI are not eliminated, but are significantly reduced, especially in the smaller background-subtraction scale maps (see Figure~20). These gains are furthered, and again significantly, by our RFI-subtraction algorithm in \textsection3.6.3. Locally modeled surfaces (\textsection1.2.1, see \textsection3.7) have been applied for visualization only. Square-root and hyperbolic-arcsine scalings are used in the top and bottom rows, respectively, to emphasize fainter structures.} \end{figure*} \begin{figure*} \plotone{fig20.pdf} \caption{\textbf{Left:} Factor by which background subtraction reduces en-route drift (red), long-duration RFI (green), large-scale astronomical signal (blue; see \textsection3.3.4), elevation-dependent signal (purple; see \textsection3.3.4), and short-duration RFI (black; see \textsection3.6.2) in our simulated data, for background subtraction scales of 3, 6, 12, and 24 beamwidths (dashed curves). Factor by which background and RFI subtraction (see \textsection3.6) reduce these contaminants (solid curves). \textbf{Right:} Fraction of the noise level to which these contaminants are reduced. If nothing is plotted, the contaminant is completely eliminated on this scale. For these measurements, each contaminant was simulated separately, and in the absence of sources.} \end{figure*} \subsection{Overview} In \textsection2, we provide a technical description of the refurbished 20-meter, including its L- and X-band receivers. We also provide a technical description of Green Bank Observatory's 40-foot telescope, used primarily for education and public engagement, but on which we developed many components of our algorithm before the 20-meter was ready. In \textsection3, we introduce our algorithm for contaminant-cleaning and mapping small-scale structures (e.g., sources) from continuum observations, and test it on both simulated and real data. In \textsection4, we demonstrate that optical-style, aperture photometry can be carried out on these maps, and introduce an algorithm for calculating photometric error bars, given the different (and more complicated, correlated) noise characteristics of these maps. We summarize our results in \textsection5. In Paper II (Dutton et al.\@ 2018): (1)~We expand on our small-scale structure algorithm to additionally contaminant-clean and map large-scale structures; (2)~We do the same for spectral (as opposed to just continuum) observations; and (3)~We carry out an X-band survey of the Galactic plane, from $-5\degr < l < 95\degr$, to showcase, and further test, techniques developed in both papers.\\ \\ \\ | In this paper, we have presented a single-dish mapping algorithm with the following features: 1. We use robust Chauvenet rejection (RCR; \textsection1.3) to improve gain calibration, making this procedure insensitive to RFI contamination (as long as it is not total, or nearly so), to catching the noise diode in transition, and to the background level ramping up or down (linearly), for whatever reason, during the calibration (\textsection3.1). 2. We again use RCR to measure the noise level of the data, in this case from point to point along the scans, also allowing this level to ramp up or down (again, linearly) over the course of the observation (\textsection3.2). We then use this noise model to background-subtract the data along each scan, without significantly biasing these data high or low (\textsection3.3). We do this by modeling the background locally, within a user-defined scale, instead of globally and hence less flexibly (as, e.g., basket-weaving approaches do). This significantly reduces, if not outright eliminates, most signal contaminants: en-route drift, long-duration (but not short-duration) RFI, astronomical signal on larger scales, and elevation-dependent signal (e.g., Figure~20). Furthermore, this procedure requires only a single mapping (also unlike basket-weaving approaches). 3. We use RCR to correct for any time delay between our signal measurements and our coordinate measurements (\textsection3.4). This method is robust against contamination by short-duration RFI and residual long-duration RFI. (In general, this procedure requires that the telescope's slew speed remain nearly constant throughout the mapping, or at least during its scans if not between them, though we do offer a modification such that it can also be applied to variable-speed, daisy mapping patterns, centered on a source.) 4. We again measure the noise level of the data, but this time from point to point across the scans, again allowing this level to ramp up or down (again, linearly) over the course of the observation (\textsection3.5). We then use this noise model to RFI-subtract the data, again without significantly biasing these data high or low (\textsection3.6). We do this by modeling the RFI-subtracted signal locally, over a user-defined scale; structures that are smaller than this scale, either along or across scans, are eliminated, including short-duration RFI, residual long-duration RFI, residual en-route drift, etc.\@ (e.g., Figure~20). This scale can be set to preserve only diffraction-limited point sources and larger structures, or it can be halved to additionally preserve Airy rings, which are visible around the brightest sources. Furthermore, this procedure can be applied to multiple observations simultaneously, in which case even smaller scales can be used (better preserving noise-level signal, and hence faint, low-S/N sources). 5. To interpolate between signal measurements, we introduce an algorithm for \textit{modeling} the data, over a user-defined weighting scale (though the algorithm can increase this scale, from place to place in the image, if more data are required for a stable, local solution; \textsection3.7). Advantages of this approach are: (1) It does not blur the image beyond its native, diffraction-limited resolution; (2) It may be applied at any stage in our contaminant-cleaning algorithm, for visualization of each step, if desired; and (3) Any pixel density may be selected. This stands in contrast to existing algorithms, which use weighted \textit{averaging} to regrid the data: (1) This does blur the image beyond its native resolution, often significantly; (2) It is usually done before contaminant cleaning takes place, because existing contaminant-cleaning algorithms -- unlike ours -- require gridded data; and (3) The pixel density is then necessarily limited to what these contaminant-cleaning algorithms can handle, computationally (\textsection1.2.3). Furthermore, since our surface-modeling algorithm does not require gridded data, images can be produced in any coordinate system, regardless of how the mapping pattern was designed. And since our surface-modeling algorithm does not assume any coordinate system-based symmetries, it works equally well with asymmetric structures. In addition to the final image, we produce a path map, a scale map, a weight map, and a correlation map, the latter three of which are important when performing photometry on the final image. 6. Lastly, we introduce an aperture-photometry algorithm for use with these images (\textsection4). In particular, we introduce a semi-empirical method for estimating photometric error bars from a single image, which is non-trivial given the non-independence of pixel values in these reconstructed images (unlike in, e.g., CCD images, where each pixel value is independent, and consequently the statistics are simpler). We also provide an empirical correction for low-S/N photometry, which can be underestimated in these reconstructed images. Additionally, we have provided a technical description of Green Bank Observatory's 20-meter diameter telescope, which we refurbished as part of this effort, and is now being used by thousands of students and researchers each year, remotely, through our Skynet interface (\textsection2.1). {\color{black}And while developed because of this telescope, the single-dish mapping and aperture-photometry algorithms presented in this paper are not specifically for this telescope. We demonstrated this by applying these algorithms to a smaller, and less capable, telescope, but there is no reason why they may not also be applied to larger, and more capable, telescopes (e.g., we recently applied them, successfully, to Green Bank Telescope mapping observations). These algorithms could also be applied to similarly collected data, with similar noise and beam characteristics, at other wavelengths, as well as from non-astronomical experiments.\footnote{\color{black}Furthermore, components of these algorithms could be used to improve interferometric reductions: E.g., RCR could be used to eliminate outlying interferometric measurements, and our surface-modeling algorithm could be modified to fill gaps in the uv plane.}} In the second paper in this series (Dutton et al.\@ 2018), we expand on the small-scale structure mapping algorithm that we have presented in this paper, in two ways: (1) We introduce an algorithm to additionally contaminant-clean and map large-scale structures, and (2) We introduce an algorithm to contaminant-clean and map spectral data, as opposed to just continuum data, as in this paper. Finally, we present an X-band survey of the Galactic plane, from $-5\degr < l < 95\degr$, to further demonstrate, and further test, the techniques that we present in both of these papers. | 18 | 8 | 1808.06128 |
1808 | 1808.03820_arXiv.txt | Powerful Laser Guide Star (LGS) systems are standard for the next-generation of extremely large telescopes. However, modern earth-based astronomy has gone through a process of concentration on few sites with exceptional sky quality, resulting in those becoming more and more crowded. The future LGS systems encounter hence an environment of surrounding astronomical installations, some of which observing with large fields-of-view. We derive formulae to calculate the impact of LGS light on the camera of a neighbouring telescope and the probabilities for a laser crossing the camera field-of-view to occur, and apply these to the specific case of the next very-high-energy gamma-ray observatory ``Cherenkov Telescope Array'' (CTA). Its southern part shall be constructed in a valley of the Cerro Armazones, Chile, close to \re{the ``Very Large Telescope'' (VLT) and} the ``European Extremely Large Telescope'' \re{(ELT)}, while its northern part will be located at the ``Observatorio del Roque de los Muchachos'', on the Canary Island of La~Palma, which also hosts the ``Gran Telescopio de Canarias'' (GTC) and serves as an optional site for the ``Thirty Meter Telescope'' (TMT), both employing LGS systems. Although finding the artificial star in the field-of-view of a CTA telescope will not disturb observations considerably, the laser beam crossing the field-of-view of a CTA telescope may be critical. We find no conflict expected for the \re{ELT} lasers, however, 1\% (3\%) of extra-galactic and 1\% (5\%) of galactic observations with the CTA may be affected by the GTC (TMT) LGS lasers, unless an enhanced version of a laser tracking control system gets implemented. | 18 | 8 | 1808.03820 |
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1808 | 1808.08135_arXiv.txt | { Solar dynamo models of Babcock-Leighton type typically assume the rise of magnetic flux tubes to be instantaneous. Solutions with high-magnetic-diffusivity have too short periods and a wrong migration of their active belts. Only the low-diffusivity regime with advective meridional flows is usually considered. } { In the present paper we discuss these assumptions and applied a time delay in the source term of the azimuthally averaged induction equation. This delay is set to be the rise time of magnetic flux tubes which supposedly form at the tachocline. We study the effect of the delay, which adds to the spacial non-locality a non-linear temporal one, in the advective but particularly in the diffusive regime. } { \cite{Fournier+17} obtained the rise time according to stellar parameters such as rotation, and the magnetic field strength at the bottom of the convection zone. These results allowed us to constrain the delay in the mean-field model used in a parameter study. } { We identify an unknown family of solutions. These solutions self-quench, and exhibit longer periods than their non-delayed counterparts. Additionally, we demonstrate that the non-linear delay is responsible for the recover of the equatorward migration of the active belts at high turbulent diffusivities. } { By introducing a non-linear temporal non-locality (the delay) in a Babcock-Leighton dynamo model, we could obtain solutions quantitatively comparable to the solar butterfly diagram in the diffusion-dominated regime. } | \label{sec:introduction} The magnetic solar cycle is attributed to a dynamo process in which motions of a conductive medium are leading to the continuous induction of magnetic fields. Those motions may arise from convection which becomes anisotropic in the presence of rotation and stratification and is approximatively described by the $\alpha$-effect \citep{KrauseRaedler80}, or from the rise of magnetic structures to the solar surface, again in the presence of rotation and stratification, leading to what is called the Babcock-Leighton effect \citep{Leighton69}. While there has been no derivation of this effect from first principles so far, the synergy with the meridional circulation in the convection zone may lead to a periodic magnetic field explaining the solar cycle satisfactorily. Babcock-Leighton dynamo were shown to reproduce qualitatively the solar butterfly diagram only if the turbulent magnetic diffusivity in the bulk of the convection zone is less than $10^{12}$~cm$^2$/s, a regime often referred to as the advective regime, since the meridional circulation then determines the cycle period. In the diffusive regime, the cycle period varies with the turbulent magnetic diffusivity and can be reduced to the observed cycle length, but the propagation of the dynamo wave then follows the Parker-Yoshimura rule and is poleward at all latitudes \citep{Yoshimura75}. A variety of Babcock-Leighton dynamos has been published during the last 25~years, exemplarily by \citet{Choudhuri+95}, \citet{Dikpati+99}, \citet{Kueker+01}, \citet{Chatterjee+04}, \citet{Guerrero+07}, and \citet{Sanchez+14} for kinematic models with stationary flows and dynamo effect, \citet{Nandy+01} for a model with toroidal-field loss by buoyancy, as well as \citet{Kitchatinov+11} for a nonlocal $\alpha$-effect similar to the Babcock-Leighton effect. A non-kinematic simulation showing a solar-like butterfly diagram is the one by \citet{Rempel06} with a turbulent magnetic diffusivity rising from $10^{10}$~cm$^2$/s to $10^{12}$~cm$^2$/s in the convection zone. Beyond the Sun, and among others, \citet{Jouve+10a} studied the cycle period dependence of Babcock-Leighton dynamos on the stellar rotation rate. The variability of the solar cycle including grand minima has been addressed with Babcock-Leighton dynamos by, for example, \citet{Karak10} who varied the meridional circulation, \citet{Olemskoy+13} who varied the strength of their nonlocal source term, and \citet{Inceoglu+17} who varied both the generation of the flows and the Babcock-Leighton effect in a non-kinematic setup. In all these attempts, however, the effect of the toroidal magnetic field in the interior of the convection zone on the poloidal-field generation at the surface is instantaneous. Note that this list of Babcock-Leighton type dynamo papers is far from complete. In the present Paper, we are addressing the problem that the Babcock-Leighton effect is not only nonlocal in space, but also in time. Following the pioneering work by \citet{Jouve+10b}, we take a further step toward a fully constrained Babcock-Leighton dynamo. Here we use the results of global numerical simulations of flux-tube rise to improve and actually constrain a Babcock-Leighton dynamo model. The global numerical simulations by \citet{Fournier+17} have shown that the rise time of magnetic flux tubes is independent of the magnetic diffusion. We suggest to treat this independence with a non-locality in space and time, by designing a Babcock-Leighton dynamo model which has a time delay in the source term. Non-localities have been shown to generate long-term variability in the amplitude and in the period of the magnetic cycle for a wide variety of models. A remarkable early attempt is the one by \citet{Yoshimura78} who used already a two-dimensional setup with source terms nonlocal both in space and time, delivering cyclic magnetic fields interrupted by ``low-activity'' periods of a few cycles duration. A sequence of papers on a zero-dimensional model, i.e. without any spatial dependence but with time delay, was published by \citet{Wilmot-Smith+06}, \citet{Hazra+14} (including stochastic variations in the Babcock-Leighton effect), and \citet{Tripathi+18}, who found sub-critical dynamo action, hinting on what we are going to show in this Paper in a two-dimensional spherical shell with realistic solar differential rotation. This is in line with the result by \citet{Rheinhardt+12} who studied the memory effect in a turbulent-$\alpha$ dynamo and found a threshold for dynamo excitation which is lower than in the case of no memory effect. There is actually a variety of papers on non-localities in turbulent-$\alpha$ dynamos which we do not review here. Closest to our approach is the work by \citet{Jouve+10b} who made the time delay magnetic-field dependent. However the authors always found that the time correlations required to obtain solar-like variability were too long to agree with the model's assumptions. Here we study such non-local models with time delays depending on the magnetic field strength, where the dependence is derived from the flux-rise simulations by \cite{Fournier+17}. Section~\ref{sec:the-model} describes the equations used and Section~\ref{sec:the-delayed-dynamos} shows the general results of using a delay in the induction equation. We discuss the implications of the results in Section~\ref{sec:discussion-and-conclusions}. | \label{sec:discussion-and-conclusions} Until now, diffusive Babcock-Leighton dynamos were considered not to be able to reproduce qualitatively the solar dynamo. The Parker-Yoshimura rule implies for the internal differential rotation of the Sun that the dynamo wave propagates poleward. We also find The cycle period to be too short as well as the low latitude radial fields to be too weak. In the present work we introduced a delay in the source term of the poloidal field. Like in \cite{Jouve+10b}, this delay represents the rise time of magnetic flux tubes through the convection zone. But in contrast to former studies, we built this model on the results of global numerical simulations of rising magnetic flux tubes in compressible stellar interiors \citep{Fournier+17}. The model consists of a rise time which depends nonlinearly on the magnetic flux density. We have shown that the nonlinearity of the delay leads to an accumulation of the Babcock-Leighton source term at certain times. When this accumulation becomes sufficiently important, it may prevent the dynamo from decaying, even though the non-delayed model shows no dynamo action. The reduction of the criticality of the dynamo opens a new window to unknown solutions. These delayed dynamos have the peculiar property of self-quenching. We found that the nonlinear delay can provide a mechanism to generate migration of the surface fields in the direction of weaker internal fields. In case of a stationary toroidal internal field at mid-latitudes, for example, the generated poloidal fields at the surface migrate towards the equator at low latitudes and towards the poles at high latitudes. This is independent of the sign of the internal differential rotation. The requirement of a low turbulent magnetic diffusivity, $\eta_{\rm t}$, for Babcock-Leighton dynamos to reproduce qualitatively the solar cycle, has been shown to be unnecessary. We demonstrate that the present delayed model, with a turbulent magnetic diffusivity of $\eta_{\rm t} = 6.7\cdot10^{11} {\rm cm}^{2}/{\rm s}$, agrees well with the solar butterfly diagram, even though the diffusivity is relatively high throughout the entire convection zone. Note that one proposed way out of the low-diffusivity problem is to use different values for $\eta_{\rm t}$ for the toroidal and for the poloidal components in the induction equation \citep{Chatterjee+04}. While the diffusivity may well be different in the horizontal and vertical directions, the poloidal field has varying components in both the horizontal and vertical directions, rendering the poloidal-field diffusivity location-dependent. We have not tried such a setup, and think it is actually not necessary given the results presented. In any case, the model presented in this work is, by design, a simplified model. It has allowed us to identify the effect of the delay on the dynamo solutions. However several ingredients are missing to reach a state-of-the-art model \citep{Rempel06, Cameron+15, Pipin17}. We only solve the induction equation for large-scale fields. We ignore the turbulent pumping, and the back-reaction of the magnetic field on the flow. All these elements will increase the complexity of the model and bring along additional free parameters which need to be constrained. The large-scale field generation based on the Babcock-Leighton effect has not been derived from first principles. Its validity remains therefore uncertain. In the absence of global simulations addressing the formation of magnetic flux tubes, the current models remain quite arbitrary. Nevertheless, the non-linearities of the presented solutions are potentially relevant for other dynamos than the Babcock-Leighton type. Finally, the relevance of this work for stellar dynamos will be revealed only if this model is proven to robustly reproduce observed dynamo patterns of further solar-like stars. | 18 | 8 | 1808.08135 |
1808 | 1808.03011_arXiv.txt | GRS~1915+105 has been active for more than 26 years since it was discovered in 1992. There are hundreds of {\it RXTE} pointed observations on this source, and the quasi-regular flares with a slow rise and a sharp decrease (i.e. the ``heartbeat" state) were recorded in more than 200 observations. The connections among the disk/corona, jet, and the disk wind at the heartbeat state have been extensively studied. In this work, we firstly perform a statistical analysis of the light curves and the X-ray spectra to investigate this peculiar state. We calculate the parameters for heartbeat cycles, including the recurrence time, the maximum and the minimum count rate, the flare amplitude, and the cumulative radiation for each cycle. The recurrence time has a bimodal distribution ranging from $\sim 20$ to $\sim 200$ s. The minimum count rate increases with increasing recurrence time; while the maximum count rate remains nearly constant around 2 Crab. Fitting the averaged spectrum for each observation, we find the strong correlations among the recurrence time, the apparent inner radius of the accretion disk (or the color correction factor), and the (nonthermal) X-ray luminosity. We suggest that the true inner edge of the accretion disk might always extend to the marginally stable orbit, while the change in corona size should result in the observed correlations. | The prototype microquasar GRS~1915+105, consisting of a black hole (BH) and a K-M \textsc{iii} companion star \citep{Greiner01}, was discovered by {\it GRANAT}/WATCH in 1992 \citep{Castro92}. Since then, the X-ray monitoring data reveal that it has remained bright during the last 26 years. The parallax distance is estimated at $\sim 8.6$ kpc, and the dynamical BH mass is measured as $\sim 12.4 M_{\odot}$ with an orbital period of 34 days \citep[][and references therein]{Reid14}. GRS~1915+105 exhibits fantastic phenomena and thus attracts a lot of attention to observations and accretion theories \citep[e.g. ][]{Fender04, Zdziarski05, McClintock06, Miller13}. Both the continuum spectral fitting and the relativistic disk reflection measurement indicate a rapidly spinning BH harbored in GRS~1915+105 \citep[e.g. ][]{Zhang97, McClintock06, Middleton06, Blum09, Miller13, Reid14}. The powerful jet with the apparent superluminal motion has been revealed \citep{Mirabel94}, and whether the jets are driven by the BH spin energy is still under dispute \citep{Fender10, Narayan12, Miller13}. The X-ray emissions show dramatic variability and quasi-periodic oscillation (QPO) in different time scales \citep[e.g. ][]{Morgan97, Belloni13, Zhang15}; therefore, GRS~1915+105 offers a remarkable laboratory for the investigation of accretion flows. Analyzing a large set of {\it RXTE} observations, \cite{Belloni00} classified 12 different variability patterns of X-ray emission, based on the count rate and the hardness ratio characteristics. Two additional variability classes were proposed later by \cite{Klein02} and \cite{Hannikainen05}. Of these 14 classes of X-ray variability, the most intriguing is the $\rho$ class, which displays the quasi-periodic bright flares, being analogous to the ``heartbeat" behavior. Therefore, the $\rho$ class is also referred to as the heartbeat state, and it is only found in two BH X-ray binaries (XRBs), i.e. GRS~1915+105 and IGR~J17091-3624 \citep[e.g. ][]{Altamirano11, Court17}. During the heartbeat state, the X-ray luminosity approaches the Eddington luminosity of the source at the peak of flares ($L_{\rm peak} \sim 10^{39}$ erg~s$^{-1}$). Thus, it provides an opportunity for us to study the accretion flows being close to the Eddington limit, which are extremely rudimentary. Both the spectral and timing properties of the heartbeat state have been extensively studied, and the connections among disk/corona, jet, and disk wind have been investigated in many papers \citep[e.g. ][]{Neilsen11, Yan17, Yan18}. But to date, there is a lack of statistical analysis on this peculiar state. In this work, we investigate the light curves and the spectra of the heartbeat state by using the full set of {\it RXTE} and search the correlations among the different parameters (Section 2). We present our results in Section 3 and explore the accretion theories at high luminosities close to the Eddington limit with the obtained relationships in Section 4. | \subsection{Movable accretion disk?} Plenty of works have been carried out to study the heartbeat state by investigating the individual observation \citep[e.g. ][]{Massaro10, Neilsen11, Neilsen12, Zoghbi16, Yan17, Yan18}. The regular flares are generally attributed to the thermal-viscous instability \citep[i.e. the Lightman-Eardley instability; ][]{Lightman74} when the inner region is dominated by radiation pressure at high luminosities \citep[e.g. ][]{Belloni97, Janiuk05, Grzedzielski17}. Phenomenologically, the burst sequence could be interpreted as the rapid emptying of the inner portion of the accretion disk, followed by a slower refilling of the inner region, and it finishes a cycle on the viscous time scale \citep{Belloni97}. Theoretically, some groups successfully reproduced the oscillating light curve with the proper cycle and the burst amplitude by modifying the classical description of $\alpha$ viscosity \citep[e.g. ][]{Honma91, Szuszkiewicz98, Merloni06, Zheng11, Grzedzielski17}. However, such models cannot account for the growth of $R_{\rm in}$ at the slow rise stage as suggested by the phase-resolved study \citep{Neilsen11}. Moreover, our statistical analysis finds the anticorrelation between $T_{\rm rec}$ and the amplitude (Figure \ref{lc_corr}), which is in contrast with the prediction of the modified viscosity model \citep{Grzedzielski17}. Besides the radiation pressure instability, \cite{Neilsen11} argued that a local Eddington effect \citep{Fukue04, Heinzeller07} is required to explain some X-ray properties of the heartbeat state. When the luminosity is larger than the critical value ($\sim 0.3-1~ L_{\rm Edd}$), the radiation pressure could either push the disk outward or generate the optically thick outflows \citep[e.g. ][]{Poutanen07, Weng14, Urquhart16}, or the disk might become thick and block the inner region \citep[e.g. ][]{McClintock06, Gu12}. As a result, the apparent size of the thermal component increases with the luminosity, and the X-ray emissions become harder in the meantime \citep{Soria07, Middleton15}. In the S branch, the X-ray luminosity is below and close to the critical luminosity, softer emissions and a smaller apparent accretion disk are observed. If the recurrence time of heartbeat flare is related to the viscous time of inner disk radius, we would expect that $T_{\rm rec}$ increases with increasing $R_{\rm in}$ and the X-ray luminosity as shown in Figure \ref{heartbeat_corr}. \subsection{Variable corona?} The advection, coronal dissipation, and outflows play important roles in the stabilization of accretion flows at the high luminosity state \citep[e.g. ][]{Janiuk02}. Investigating the quasi-simultaneous radio and X-ray data, \cite{Vadawale03} achieved the association between jet and corona. For instance, the central Compton cloud (corona) is ejected during the soft X-ray dips, which are preceded by a radio-loud hard state. They argued that the change in the corona can account for the X-ray variations that were previously attributed to the accretion disk. Such a scenario, in particular, is supported by the absence of both the Comptionized component and low-frequency QPOs during the X-ray dips. Note that low-frequency QPO signals detected in low-mass X-ray binaries are generally connected with the hard component but not (or sometimes indirectly linked with) the thermal accretion disk \citep[e.g. ][]{Belloni14, Motta15, Zhao16, Zhang17, Yan18}. Our results indicate that, for observations with different $T_{\rm rec}$, the main difference is from the slow rise stage, while the source takes almost the same time to return from the peak luminosity to the minimum (Figure \ref{lc}). The change of $C_{\rm min}$ would further point to the variation of corona and jet properties for different $T_{\rm rec}$ oscillations. The tighter correlation between $R_{\rm in}$ and the power-law luminosity (Figure \ref{heartbeat_corr}) also supports that the limit cycle is driven by the nonthermal emissions rather than the thermal component. Meanwhile, $C_{\rm max}$ remains nearly constant, indicating that the bolometric luminosity of GRS~1915+105 reaches its Eddington luminosity. The variation of the corona does not only affect the Comptionized component itself, but also modifies the property of the observed disk thermal emission via Compton scattering \citep[e.g. ][]{Shimura95, Nayakshin00}. The spectral hardening factor $f$ is by no means constant, as we assumed in the $R_{\rm in}$ calculation. It should depend on the accretion rate, the fraction of the accreted power released in the corona, the corona reflecting, etc \citep[e.g. ][]{Merloni00, Davis05}. When the actual inner boundary of the accretion disk is fixed, the color disk radius could vary by more than a factor of 4 with different accretion parameters \citep{Merloni00}. Additionally, \cite{Zoghbi16} argued that a change in composition of disk atmosphere could make an even larger change in the color correction factor. As the disk temperature and flux increase, more iron ions are elevated up by the radiation force, increasing the opacity in the upper disk and resulting in smaller $f$. They further suggested that the evolution of $R_{\rm in}$ inferred from the disk blackbody component was artificial, while the actual inner edge was kept at a small radius ($\sim 1.1~R_{\rm g}$) in the heartbeat state according to the reflection measurement. Specifically, \cite{Zoghbi16} proposed that the inner disk radius remained constant and the corona changed size during oscillations. If the corona is smaller and closer to the central BH, more photons are dragged by the strong gravity and hit the disk; therefore, fewer photons can reach the observer. This model is in qualitative agreement with our results presented in Figure \ref{comptt}, i.e. the {\it comptt}-to-total-flux ratio and $T_{\rm rec}$ increase with the total X-ray luminosity. The change of apparent disk size shown in Figure \ref{spec_corr} can be interpreted as the variation of the color correction factor. The correlation between the X-ray/disk flux and $Norm_{\rm diskpn}$ (Figure \ref{comptt}) indicates that the factor $f$ decreases from $\sim 3.0$ to $\sim 2.4$ with the increasing luminosity. But note that if the inner disk radius was fixed to 6 $R_{\rm g}$, the actual value of $f$ would be smaller with smaller inner disk radius. Although the varying corona scenario offers a promising explanation for the results given by our statistical analysis, we cannot completely rule out the evolving disk model. In addition, a number of questions remain to be answered. (1) What is the origin of variations in the corona (and the disk)? (2) If the low-frequency QPOs and $T_{\rm rec}$ are account for different time scales of corona, e.g. dynamical time and thermal/viscous time; but their correlation is unknown at the current stage. (3) It has been pointed out that the reflection component is prominent in the heartbeat state \citep{Zoghbi16}. The evolution trend of the disk component and the total X-ray luminosity are insensitive to models with or without the reflection. However, the parameters of the thermal and Comptonization components might change when the reflection is included or excluded in fitting models. Unfortunately, we cannot study the reflection spectrum with the {\it RXTE}/PCA data due to its low energy resolution. On the other hand, the single {\it Nustar} observation is inadequate to explore the spectral/temporal evolution with different $T_{\rm rec}$. Thus, we would like to suggest that multiple high quality data (high-energy/time resolution) from Insight-{\it HXMT} and {\it Nustar} are required to probe the peculiar state deeper. | 18 | 8 | 1808.03011 |
1808 | 1808.08904_arXiv.txt | Recent numerical work in helioseismology has shown that a periodically varying subsurface magnetic field leads to a fanning of the $f$-mode, which emerges from a density jump at the surface. In an attempt to model a more realistic situation, we now modulate this periodic variation with an envelope, giving thus more emphasis on localised bipolar magnetic structures in the middle of the domain. Some notable findings are: (i) compared to the purely hydrodynamic case, the strength of the $f$-mode is significantly larger at high horizontal wavenumbers $k$, but the fanning is weaker for the localised subsurface magnetic field concentrations investigated here than the periodic ones studied earlier; (ii) when the strength of the magnetic field is enhanced at a fixed depth below the surface, the fanning of the $f$-mode in the $k\omega$ diagram increases proportionally in such a way that the normalised $f$-mode strengths remain nearly the same in different such cases; (iii) the unstable Bloch modes reported previously in case of harmonically varying magnetic fields are now completely absent when more realistic localised magnetic field concentrations are imposed beneath the surface, thus suggesting that the Bloch modes are unlikely to be supported during most phases of the solar cycle; (iv) the $f$-mode strength appears to depend also on the depth of magnetic field concentrations such that it shows a relative decrement when the maximum of the magnetic field is moved to a deeper layer. We argue that detections of $f$-mode perturbations such as those being explored here could be effective tracers of solar magnetic fields below the photosphere before these are directly detectable as visible manifestations in terms of active regions or sunspots. | The Sun supports a wide variety of waves that carry useful information about the internal solar structure, which can be inferred by employing the methods of helioseismology \citep[see, e.g.,][]{G87,C03}. Local analysis using especially the surface gravity mode (or the $f$-mode) is useful in studying the near-surface structure \citep[e.g.][]{HBBG08,FBCB12,FCB13,DABCG11}. Some properties of surface waves in idealised settings involving magnetic fields were explored in \citet{C61,R81,MR89,MR92,MAR92}. It is of great interest to probe interior magnetic fields of the Sun using the techniques of helioseismology; see \cite{T06} for a review on the subject of magnetohelioseismology. Systematic changes, notably in the frequencies of the global $f$-mode, as a function of solar cycle were found and discussed in detail in \citet{T06} and \citet{P08}. \cite{T06} found evidence of a $500\G$ magnetic field at a depth of about $5\Mm$ and suggested $2\%$ modulation in turbulent convection velocities from solar minimum to maximum, based on the observed cycle-dependence of mean frequency shifts of the $f$- and $p$-modes. Moreover, it is reasonable to expect magnetically induced variations in the $f$-mode on much shorter timescales, e.g., during localised magnetic flux emergence leading to the formation of active regions (ARs) or sunspots. Much of the earlier studies on the global $f$-mode of the Sun focussed primarily on the frequency shifts that were observed \citep{LWK90,FSTT92}. The frequencies were significantly smaller than the theoretically expected values, where both the shift and line width grow with the spherical harmonic degree. Subsequent studies explained these findings by invoking turbulent background motions and deriving a generalised dispersion relation of the $f$-mode in the presence of a random velocity field \citep{MR93a,MR93b,MMR99,M00a,M00b,MKE08}. The influence of coherent (for modelling supergranulation) as well as random (mimicking near-surface granulation) flows on the $f$-mode were explored in detail by \cite{M00a} where it was found that, while a space-dependent random flow causes a decrement, a time-dependent random flow can enhance the frequencies. Observations were thus explained in terms of the parameters of the chosen velocity field. However, these studies ignored the effect of magnetic fields, which can increase the $f$-mode frequencies \citep[e.g.][]{C61,R81,MR92,MAR92}. In a series of works, \citet{CB93,CB97} and \citet{CBZ94} investigated the interaction of $f$- and $p$-modes with a vertical magnetic field and found that a partial conversion of these modes into slow magnetoacoustic modes takes place whenever they encounter a vertical field resembling those of sunspots. \cite{PK09} numerically explored the effects of inclined magnetic fields and noted that the $f$-modes are more strongly affected by the background magnetic field than the $p$-modes. \cite{SBCR15} studied numerically various properties of the $f$- as well as $p$- and $g$-modes in a wide variety of magnetic backgrounds. They found that horizontal magnetic fields cause an increase in the $f$-mode frequencies -- as expected. But their dependencies are more complicated in the presence of vertical or oblique magnetic fields, which may be more relevant for the predominantly vertical fields of sunspots. In this case, the $f$-mode frequencies are enhanced relative to their nonmagnetic values at intermediate horizontal wavenumbers, but decreased at large wavenumbers. In the presence of magnetic fields in the solar atmosphere, Alfv\'en and magnetosonic waves are known to couple resonantly with the global oscillations, affecting mainly the frequencies, line widths, and penetration of the $f$ and $p$ modes into the solar atmosphere \citep[see][for a review and references therein]{E06}. Another important finding was that resonant interaction between global modes and Alfv\'en waves causes a damping of the $f$- and $p$-modes due to dissipative effects near the resonance frequency. \cite{PE18} further study this situation by deriving the dispersion relation of the $f$-mode in a magnetically coupled solar interior-atmosphere system to obtain its frequency shifts. For a magnetised atmosphere, these were found to be positive relative to an unmagnetised one. It is of great interest to use numerical simulations to study the effects of subsurface magnetic fields on both the acoustic or $p$-modes and the $f$-mode. Such studies aim at refining the interpretation of helioseismic measurements which use sound waves to infer the internal structure of the Sun and its internal motions \citep[e.g.][]{Ba16,HGR16}. This is necessary, because magnetic fields complicate the usage of helioseismic inversion techniques, as their presence gives rise to the modification of sound and gravity waves into magnetoacoustic and magnetogravity ones which are difficult to account for \citep{T83,C11}. Furthermore, if the properties and behaviour of $p$- and $f$-modes in the presence of magnetic fields of different strength, topology or location, in particular depth, were known well enough from the numerical models, we might be able to infer the subsurface magnetic fields from helioseismic measurements. From such inferences, we may be able to learn about the origin of subsurface magnetic fields, that is, about the solar dynamo mechanism \citep{B05,Cha10,KMB12}, and about the process of concentrating magnetic fields into sunspots and ARs; see, e.g., \cite{BRK16} and \cite{KBKKR16} for a competing mechanism enabling a self-consistent formation of localised magnetic flux concentrations. Such numerical simulations, enabling us to study the effects of magnetic fields on the naturally occurring modes of oscillations, have recently been performed with the {\sc Pencil Code}\footnote{http://github.com/pencil-code}. \cite{SBCR15} introduced a modelling framework with a piecewise isothermal atmosphere, where the upper layer is mimicking a hot corona and the lower one a (cooler) convection zone. The two layers are separated by a jump in density and temperature, which represents the solar surface and enables the presence of the $f$-mode. For a simple approximation of the convective turbulence, random hydrodynamic forcing was applied in the lower layer. In such a setup, acoustic ($p$), internal gravity ($g$), and surface gravity ($f$) modes are all {\it self-consistently} driven, in contrast to other types of approaches, which selectively produce the modes based on linearised equations \citep{DABCG11,SCGM11}. \cite{SBCR15} studied the influence of uniformly imposed magnetic fields on the $p$- and $f$-mode properties, and verified the expectation of $f$-modes being sensitive to the presence of magnetic fields and $p$-modes being less affected. The major effect on the $f$-modes was an increase in their mode frequencies. \cite{SBR14} developed the setup further to include non-uniform magnetic fields with harmonic profiles within the convection zone. This work revealed the {\it fanning effect} of the $f$-mode, that is, an increase of its line width towards higher wavenumbers. The width of the fan and its asymmetry could be directly related to the strength and location of the magnetic field. These numerical studies and their predictions led to an observational case study of $f$-modes in relation to the emergence of about half a dozen ARs, observed with HMI, and the resulting line-of-sight Dopplergrams and magnetograms \citep{SRB16}. This study reported \emph{strengthening} of the local $f$-mode about two days before the emergence of ARs at the same corotating patch. It was noted that such a precursor signal can be detected by isolating high-degree $f$-modes through careful fitting and subsequent subtraction of background and $p$-modes. It was argued there that the precursor signal is best seen when an AR forms in isolation, i.e., far from other existing ARs which can `pollute' the signal. It is known that the sunspots or ARs absorb the $f$-mode power, thus causing its damping \citep{CB97}, and therefore it is expected to be harder to extract the precursor signal associated with a newly forming AR in a `crowded' environment with many existing ARs. A possible cleaning procedure to still extract the signal was discussed in \cite{SRB16}. Although this is yet to be confirmed for larger data sets and with independent observational techniques, the potential significance of such $f$-mode related precursors cannot be understated. Based on the photospheric velocity measurements, \cite{KT17} detected plasma upflows somewhat before the emergence of two small ARs. A number of previous case studies have reported detections of subphotospheric velocities associated with emerging ARs using different techniques, such as, helioseismic holography, ring-diagram or time-distance analysis \citep{Komm08,Har11,Ilo11,Birch13,Barnes14}. Given that the $f$-mode eigenfunctions extend to depths of about a few Mm, it might be expected to be sensitive to such velocity perturbations. In this paper we extend the model of \cite{SBR14,SBCR15} to study the $f$-mode strengthening and fanning using inhomogeneous magnetic fields of different strength, topology, and location in the convection zone. In addition to using harmonic profiles which stretch over the whole horizontal extent, we use localised harmonic perturbations, to better mimic isolated active regions. We describe our model and basic definitions in \Sec{model} and present our results in \Sec{res}. Conclusions are given in \Sec{conc}. | \label{conc} We have shown numerically that the $f$-mode is significantly perturbed in the presence of subsurface magnetic fields concentrated one to a few scale heights below the photosphere. The fanning (or increase in line width) of the $f$-mode was first reported in \cite{SBR14} based on harmonically varying magnetic fields, but the associated mode strengthening was not emphasised there. Here we extended this by also investigating the effects of localised bipolar magnetic fields, resembling more realistically the active region precursors. The fanning effect of the $f$-mode is found to be weaker in the case of localised magnetic field concentrations. Motivated by the observational findings of \cite{SRB16}, where the strengthening at large wavenumbers of the $f$-mode prior to the emergence of about half a dozen active regions, we focus here primarily on the phenomenon of $f$-mode strengthening. In our numerical investigations reported here, the $f$-mode strength is found to be clearly larger at high horizontal wavenumbers in a variety of magnetic background states, compared to their corresponding hydrodynamic cases. While the fanning of the $f$-mode at large $k_x$ is more sensitive to the strength of the magnetic field at a given depth, the mode strength itself shows a dependence on the location of the magnetic concentration beneath the surface. Thus, it is indeed remarkable that the properties of the $f$-mode trace so effectively both the location and the strength of subsurface magnetic fields that are not yet directly visible at the photosphere. We note that damping of the $f$-mode due to its resonant coupling with atmospheric magnetic fields or due to absorption caused by sunspots is known \citep{CB97,E06,PEG07}, and also seen in observational studies of \cite{SRB16} \emph{after} the emergence of active regions. But, the strengthening of the $f$-mode due to magnetic fields confined below the photosphere is a qualitatively different finding, which is numerically confirmed in the present work. Let us now estimate the instrumental requirements to robustly probe and detect solar $f$-mode perturbations of the kind explored in this paper. In order to detect both fanning and mode-strength gain, according to \Figs{fmm_A}{fmm_BI}, $\widetilde{k}_x\gtrsim2$ needs to be reached. With $\gamma \Hd \approx 0.5\Mm$ and $R_\odot=700\Mm$ being the solar radius, this corresponds to spherical harmonic degrees $\ell \gtrsim 2800$. However, indications of $f$-mode strengthening prior to active region formation could indeed be available even at somewhat smaller degrees, as was seen in the observational work of \cite{SRB16} based on data from the HMI instrument, whose sensitivity drops beyond $\ell\approx2000$. The precursor signal is expected to improve at larger wavenumbers and therefore higher resolution data from other facilities and upcoming missions will be critical not only for establishing the connection between the solar $f$-mode and the subsurface magnetic fields, but also for providing deeper insights into the sunspot formation mechanism, which is still an outstanding topic in solar physics. More work is needed to understand the physical mechanism responsible for such a strengthening of the $f$-mode and this will be attempted elsewhere. In particular, it would be useful to extend these models to explore effects of 3-dimensional magnetic configurations; see the review of \cite{E06} for possible effects on mode damping and mode conversion. | 18 | 8 | 1808.08904 |
1808 | 1808.06634_arXiv.txt | Gravitational potential fluctuations driven by bursty star formation can kinematically `heat up' dark matter at the centres of dwarf galaxies. A key prediction of such models is that, at a fixed dark matter halo mass, dwarfs with a higher stellar mass will have a lower central dark matter density. We use stellar kinematics and HI gas rotation curves to infer the inner dark matter densities of eight dwarf spheroidal and eight dwarf irregular galaxies with a wide range of star formation histories. For all galaxies, we estimate the dark matter density at a common radius of 150\,pc, $\rho_{\rm DM}(150\,\mathrm{pc})$. We find that our sample of dwarfs falls into two distinct classes. Those that stopped forming stars over $6$\,Gyrs ago favour central densities $\rho_{\rm DM}(150\,\mathrm{pc})>10^8$\,M$_{\odot}$\,kpc$^{-3}$, consistent with cold dark matter cusps, while those with more extended star formation favour $\rho_{\rm DM}(150\,\mathrm{pc})<10^8$\,M$_{\odot}$\,kpc$^{-3}$, consistent with shallower dark matter cores. Using abundance matching to infer pre-infall halo masses, $M_{200}$, we show that this dichotomy is in excellent agreement with models in which dark matter is heated up by bursty star formation. In particular, we find that $\rho_{\rm DM}(150\,\mathrm{pc})$ steadily decreases with increasing stellar mass-to-halo mass ratio, $M_*/M_{200}$. Our results suggest that, to leading order, dark matter is a cold, collisionless, fluid that can be kinematically `heated up' and moved around. | \label{sec:intro} The standard $\Lambda$ Cold Dark Matter ($\Lambda$CDM) cosmological model gives a remarkable description of the growth of structure in the Universe on large scales \citep[e.g.][]{2006Natur.440.1137S,2006ApJ...648L.109C,2013AJ....145...10D,2014A&A...571A..16P,2016JCAP...08..012B,2016MNRAS.456.2301W}. Yet, on smaller scales inside the dark matter halos of galaxies, there have been long-standing tensions \citep[e.g.][]{2017ARA&A..55..343B}. The oldest of these is the `cusp-core' problem. Pure dark matter (DM) structure formation simulations in $\Lambda$CDM predict a universal DM halo profile that has a dense `cusp' at the centre, with inner density $\rho_{\rm DM} \propto r^{-1}$ \citep{1991ApJ...378..496D,1996ApJ...462..563N}. By contrast, observations of gas rich dwarf galaxy rotation curves have long favoured DM `cores', with $\rho_{\rm DM} \sim {\rm constant}$ \citep{1994ApJ...427L...1F,1994Natur.370..629M,2010AdAst2010E...5D,2017MNRAS.467.2019R}. The cusp-core problem has generated substantial interest over the past two decades because it may point to physics beyond the collisionless `Cold Dark Matter' (CDM) typically assumed to date. \citet{2000PhRvL..84.3760S} were the first to suggest that `Self Interacting Dark Matter' (SIDM) -- that invokes a new force acting purely in the dark sector -- could transform a dense cusp to a core through energy transfer between the DM particles \citep[e.g.][]{2013MNRAS.430...81R,2015MNRAS.453...29E,2016PhRvL.116d1302K,2017MNRAS.470.1542S,2017MNRAS.472.2945R}. Warm Dark Matter (WDM) has also been proposed as a solution to the cusp-core problem (e.g. \citealt{2000PhRvD..62f3511H,2001ApJ...556...93B,2001ApJ...559..516A,2014MNRAS.439..300L,2017MNRAS.470.1542S}, but see \citealt{2001ApJ...561...35D,2012MNRAS.424.1105M} and \citealt{2013MNRAS.430.2346S}). Other solutions include `fuzzy DM' \citep{2000PhRvL..85.1158H,2017PhRvD..95d3541H}, `fluid' DM \citep{2000ApJ...534L.127P} and `wave-like' DM \citep{2014NatPh..10..496S}. However, there is a more prosaic explanation for the cusp-core problem. If gas is slowly accreted onto a dwarf galaxy and then suddenly removed (for example by stellar winds or supernovae feedback) this causes the DM halo to expand, irreversibly lowering its central density\footnote{Note that there is an alternative mechanism by which stars and gas can alter the inner DM density profile. \citet{2001ApJ...560..636E} were the first to suggest that dense gas clumps could impart angular momentum to the inner DM density profile by dynamical friction, causing a cusp to flatten to a core (see also \citealt{2009ApJ...698.2093D}, \citealt{2010ApJ...725.1707G} and \citealt{2011MNRAS.416.1118C} for more recent work on this). Such a mechanism still requires stellar feedback to then destroy these dense gas clumps. Otherwise, the inner stellar density that results would be too high to be consistent with observations \citep[e.g.][]{2015MNRAS.446.1820N}. The predictions from this class of model can be rather degenerate with `DM heating' due to potential fluctuations \citep{2016Ap&SS.361..162D} and it may well be that both act in tandem in dwarf galaxies. This remains an area of active research.} \citep{1996MNRAS.283L..72N}. \citet{2002MNRAS.333..299G} showed that, for reasonable gas fractions and collapse factors, the overall effect of this `DM heating' is small. However, if the effect repeats over several cycles of star formation, it accumulates, leading eventually to complete DM core formation \citep{2005MNRAS.356..107R}. Indeed, recent numerical simulations of dwarf galaxies that resolve the impact of individual supernovae on the interstellar medium find that the gas mass within the projected half light radius of the stars, $R_{1/2}$, naturally rises and falls on a timescale comparable to the local dynamical time\footnote{Fluctuations in the central gas mass need not be very large to excite DM heating, so long as they are repeated enough times. \citet{2016MNRAS.459.2573R} find in their simulations that the dynamical mass within $R_{1/2}$ fluctuates by just $\sim 10$\%, yet this is sufficient transform a DM cusp to a core within $R_{1/2}$.}, transforming an initial DM cusp to a core (e.g. \citealt{2008Sci...319..174M,2012MNRAS.421.3464P,2015MNRAS.454.2092O,2016MNRAS.456.3542T,2016MNRAS.459.2573R}, and for a review see \citealt{2014Natur.506..171P}). Such simulations have already made several testable predictions. \citet{2013MNRAS.429.3068T} show that the gas flows that transform DM cusps to cores lead to a bursty star formation history, with a peak-to-trough ratio of $5-10$ and a duty cycle comparable to the local dynamical time. Furthermore, the stars are dynamically `heated' similarly to the DM, leading to a stellar velocity dispersion that approaches the local rotational velocity of the stars ($v/\sigma \sim 1$) inside $R_{1/2}$. Both of these predictions are supported by observations of dwarf galaxies \citep[e.g.][]{1998AJ....116.1227D,2002AJ....123..813D,2007ApJ...659..331Y,2014MNRAS.441.2717K,2012ApJ...750...33L,2017MNRAS.465.2420W,2017MNRAS.466...88S}. Further evidences for `DM heating' come from the observed age gradients in dwarfs \citep{2016ApJ...820..131E}. While there is strong evidence that dwarf galaxies have bursty star formation histories, this is only circumstantial evidence for DM heating. The real `smoking gun' for DM cusp-core transformations lies in another key prediction from recent numerical models: DM core formation requires several cycles of gas inflow and outflow \citep{2005MNRAS.356..107R,2012MNRAS.421.3464P}. Thus, at fixed halo mass, galaxies that have formed more stars (i.e. that have undergone more gas inflow-outflow cycles) will have a lower central DM density \citep{2012MNRAS.421.3464P,2012ApJ...759L..42P,2014MNRAS.437..415D,2015MNRAS.454.2092O,2015MNRAS.450.3920B,2016MNRAS.459.2573R,2017MNRAS.466L...1D,2018arXiv180607679B}. By contrast, solutions to the cusp-core problem that invoke exotic DM predict no relationship between the central DM densities of dwarfs and their star formation histories\footnote{Most exotic DM models designed to solve the cusp-core problem predict that {\it all} dwarfs should have a central DM core. However, there can be exceptions to this. In SIDM models with a high self-interaction cross section, for example, dark matter halos undergo `core collapse', increasing their central density at late times \citep[e.g.][]{2012MNRAS.423.3740V}. However, while this will lead to some stochasticity in the central DM density of dwarfs, it will not lead to any relationship between their central DM densities and their star formation histories.}. Whether or not a dwarf will form a DM core depends primarily on the number and amplitude of gas inflow-outflow cycles, and on the amount of DM that needs to be evacuated from the centre of the dwarf to form the core. This can be posed in the form of an energy argument, whereby the total energy available to move gas around depends on the total stellar mass formed, $M_*$, while the energy required to unbind the DM cusp depends on the DM halo mass, $M_{200}$ \citep{2012ApJ...759L..42P}. Thus, whether or not a DM core will form in a given dwarf galaxy depends primarily on its stellar mass to halo mass {\it ratio}, $M_*/M_{200}$ \citep{2014MNRAS.437..415D}. However, since $M_{200}$ is challenging to extrapolate from the data, in this paper we consider also a proxy for the ratio $M_*/M_{200}$: the star formation `truncation time', $t_{\rm trunc}$. We define this to be the time when the dwarf's star formation rate (SFR) fell by a factor of two from its peak value\footnote{This is similar to the concept of `fast' and `slow' dwarfs introduced by \citet{2015ApJ...811L..18G} and explored in more detail by \citet{2018arXiv180607679B}. However, our definition here is more readily applied to our sample of both dSphs and dIrrs (see also \citet{2018arXiv180707093R} for a discussion on this point).}. This can be used as a proxy for $M_*/M_{200}$ so long as the SFR is approximately constant\footnote{If dwarfs have significant gaps in their star formation histories, then this correspondence between $t_{\rm trunc}$ and $M_*/M_{200}$ can break \citep[e.g.][]{2019MNRAS.482.1176W}. For this reason, in this paper we will look for anti-correlations between the central DM density of dwarfs and $t_{\rm trunc}$ (that is easier to measure) and $M_*/M_{200}$ (that is more fundamental, but harder to estimate).} (as is the case for the sample of dwarfs that we consider in this paper; see \citealt{2018arXiv180707093R} and \S\ref{sec:data}). In this case, dwarfs with $t_{\rm trunc} \rightarrow 0$\,Gyrs have $M_*/M_{200} \rightarrow 0$, while those with $t_{\rm trunc} \rightarrow 13.8$\,Gyrs (i.e. the age of the Universe) have formed stars for as long as possible and have, therefore, maximised both $M_*/M_{200}$ and their ability to produce a DM core. Unlike $M_{200}$, however, $t_{\rm trunc}$ has the advantage that it is readily estimated from the dwarf's star formation history (SFH; see \S\ref{sec:data}). In this paper, we set out to test the above key prediction of DM heating models, that dwarfs with `extended star formation' (i.e. $t_{\rm trunc} \rightarrow 13.8$\,Gyrs and maximal $M_*/M_{200}$) have DM cores, while those with `truncated star formation' (i.e. $t_{\rm trunc} \rightarrow 0$\,Gyrs and minimal $M_*/M_{200}$) have DM cusps. To achieve this, we estimate the central DM density, $M_*$, $t_{\rm trunc}$ and $M_{200}$ for a sample of nearby dwarf galaxies with a wide range of star formation histories (SFHs). Our sample includes gas-poor dwarf spheroidal galaxies (dSphs) whose star formation ceased shortly after the beginning of the Universe, dSphs with extended star formation that shut down only very recently, and gas rich dwarf irregular galaxies (dIrrs) that are still forming stars today. This requires us to accurately infer the DM distribution in both gas rich and gas poor galaxies. For the former, we use HI rotation curves as in \citet{2017MNRAS.467.2019R}; for the latter, we use line of sight stellar kinematics. However, with only line of sight velocities, there is a well-known degeneracy between the radial density profile (that we would like to measure) and the velocity anisotropy of the dwarf (see \S\ref{sec:gravsphere} and \citealt{1982MNRAS.200..361B,1990AJ.....99.1548M,2013NewAR..57...52B,2017MNRAS.471.4541R}). In \citet{2017MNRAS.471.4541R} and \citet{Read:2018pft}, we introduced a new mass modelling tool -- \GravSphere\ -- that breaks this degeneracy by using `Virial Shape Parameters' (VSPs). We used a large suite of mock data to demonstrate that with $\sim 500$ radial velocities, \GravSphere\ is able to correctly infer the dark matter density profile over the radial range $0.5 < r/R_{1/2} < 2$, within its 95\% confidence intervals. Here, we use \GravSphere\ to infer the inner DM density of eight Milky Way dSphs that have radial velocities for $\simgt 190$ member stars. We emphasise that, while with of order 500 radial velocities, \GravSphere\ is not able to obtain a robust inference of the inner {\it slope} of the DM density profile, it can constrain the {\it amplitude} of the inner DM density at $\sim 150$\,pc \citep{Read:2018pft}. As we shall show, this is sufficient to test DM heating models. This paper is organised as follows. In \S\ref{sec:cuspcoretheory}, we briefly review the cusp-core problem in $\Lambda$CDM, and we explain why a robust estimate of the amplitude of the DM density at 150\,pc is sufficient for testing DM heating models. In \S\ref{sec:method}, we describe our method for measuring the DM density profile from stellar kinematics (\GravSphere; \S\ref{sec:gravsphere}) and HI rotation curves (\S\ref{sec:HIrot}). In \S\ref{sec:data}, we describe our data compilation for our sample of dIrrs and dSphs, including their SFHs and estimates of $M_{200}$ taken from the literature. In \S\ref{sec:results}, we present our key results. In \S\ref{sec:discussion}, we compare our measurements with previous work in the literature. We discuss the robustness of our results and their implications for `DM heating' and the nature of DM. Finally, in \S\ref{sec:conclusions} we present our conclusions. | \label{sec:conclusions} We have used stellar kinematics and HI rotation curves to infer the radial DM density profile of eight dwarf spheroidal (dSph) and eight dwarf irregular (dIrr) galaxies with a wide range of star formation histories. Our key findings are as follows: \begin{itemize} \item The dwarfs fell into two distinct classes. Galaxies with only old stars ($>6$\,Gyrs old) had central DM densities, $\rho_{\rm DM}(150\,{\rm pc}) > 10^8$\,M$_\odot$\,kpc$^{-3}$, consistent with DM cusps; those with star formation until at least 3\,Gyrs ago had $\rho_{\rm DM}(150\,{\rm pc}) < 10^8$\,M$_\odot$\,kpc$^{-3}$, consistent with DM cores (Figure \ref{fig:dmheating}, left panel). \item We estimated pre-infall halo masses for our sample of dwarfs, using HI rotation curve measurements for the dIrr sample and abundance matching for the dSph sample. With this, we showed that their $\rho_{\rm DM}(150\,{\rm pc})$ as a function of $M_{200}$ is in good agreement with models in which DM is kinematically `heated up' by bursty star formation. The dwarfs with only old-age stars lay along the track predicted by the NFW profile in $\Lambda$CDM, consistent with having undergone no measurable DM heating. By contrast, those with extended star formation lay along the track predicted by the \coreNFW\ profile from \citet{2016MNRAS.459.2573R}, consistent with maximal DM heating (Figure \ref{fig:dmheating}, middle panel). \item We found that $\rho_{\rm DM}(150\,{\rm pc})$ for our sample of dwarfs is anti-correlated with their stellar mass to pre-infall halo mass ratio, $M_*/M_{200}$ (Figure \ref{fig:dmheating}, right panel). This is also in good quantitative agreement with predictions from recent DM heating models \citep{2012ApJ...759L..42P,2014MNRAS.437..415D,2015MNRAS.454.2981C,2016MNRAS.459.2573R,2016MNRAS.456.3542T}. \item While not the main focus on this paper, in Appendix \ref{app:dwarftwins} we discussed the implications of our results for `alternative gravity' models for DM. There, we showed that the dwarf `twins' Draco and Carina provide a particularly clean test of such models. These two dwarfs have similar $M_*$, $R_{1/2}$ and orbit around the Milky Way, yet favour very different dark matter density profiles. In $\Lambda$CDM, this is explained by Carina and Draco inhabiting halos with different pre-infall masses and concentrations (Figure \ref{fig:dmheating}, middle panel). In alternative gravity theories, however, the existence of visibly similar galaxies with different gravitational force-fields represents a major challenge (Figure \ref{fig:mond}). \end{itemize} | 18 | 8 | 1808.06634 |
1808 | 1808.08935.txt | \baselineskip=15pt Among a number of active galactic nuclei (AGNs) that drive ionized outflows in X-rays, a low-redshift ($z = 0.184$) quasar, PDS~456, is long known to exhibit one of the exemplary ultra-fast outflows (UFOs). %It has been extensively analyzed over many years. However, the physical process of acceleration mechanism is yet to be definitively {constrained}. In this work, we model the variations of the Fe K UFO properties in \obj\ over many epochs in X-ray observations %in 2013/2014 {\it XMM-Newton}/{\it NuSTAR} spectra in the context of magnetohydrodynamic (MHD) accretion-disk winds employed in our earlier studies of similar X-ray absorbers. % We applied the model to the 2013/2014 {\it XMM-Newton}/{\it NuSTAR} spectra to determine the UFO's condition; namely, velocity, ionization parameter, column density and equivalent width (EW). % %The model further account for the reported {variations} of the detected UFO's equivalent width (EW) and velocity on the X-ray . Under some provisions on the dependence of X-ray luminosity on the accretion rate applicable to near-Eddington state, our photoionization calculations, coupled to a 2.5-dimensional MHD-driven wind model, can further reproduce the observed correlations of the UFO velocity and the anti-correlation of its EW with X-ray strength of \obj. %We find that the competition between X-ray strength and plasma wind density in the massive wind can naturally lead to the variable nature of the UFOs consistent with data for \obj. % This work demonstrates that UFOs, even without radiative pressure, can be driven as an extreme case purely by magnetic interaction while also producing the observed spectrum and correlations. | One of the generic features seen in black hole (BH) systems such as active galactic nuclei (AGNs) and Galactic X-ray binaries (XRBs) are blue-shifted absorption features in their spectra primarily detected in the UV and the X-ray bands, with the latter also known as warm absorbers (WAs). % %In addition, it has become clearer, owing to high-resolution data in recent years, that more than $50\%$ of the Seyfert 1 AGNs seem to exhibit such X-ray-absorbing outflows -- ionized X-ray winds also known as warm absorbers (WAs) -- being ejected at typical velocities ranging from hundreds to thousands of km~s$^{-1}$ with characteristic column densities of $N_H \sim 10^{20-22}$ cm$^{-2}$ over a wide range of ionization parameter defined by $\xi \equiv L_{\rm ion} / (n r^2)$ where $L_{\rm ion}$ is ionizing (X-ray) luminosity (likely to originate from a compact central region such as coronae), $n$ is plasma number density at distance $r$ from the BH. While historically the discovery of the WAs dates back to the {\it ASCA} era when the bound-free edge features from \ovii/\oviii\ were definitively found for the first time as a consequence of photoionization, it has been well known to date that various ions of different charge state are also outflowing. A more detailed spectral analyses have been made possible with grating spectroscopy of high-resolving power as in {\it Chandra}/HETG and {\it XMM-Newton}/RGS... % Within the last decade, on the other hand, a new class of outflows has drawn much attention because of their unique physical characteristics: They are ejected at near-relativistic velocities ($v/c \sim 0.1-0.7$), with nearly Compton-thick columns ($10^{23} \lsim N_H \lsim 10^{24}$ cm$^{-2}$) and systematically high ionization parameter\footnote[1]{This is defined as $\xi \equiv L_{\rm ion} / (n r^2)$ where $L_{\rm ion}$ is ionizing (X-ray) luminosity, $n$ is plasma number density at distance $r$ from the BH.} ($\log \xi \sim 4-6$). These ultra-fast outflows (UFOs), primarily observed with high-throughput CCD detectors, appear to be present not only in nearby Seyfert AGNs \citep[e.g.][]{Tombesi10,Tombesi11,Tombesi14} but also in very bright (lensed) quasars \citep[e.g.][]{Chartas03,Pounds03,Chartas09a,Dadina18} and presumably in the ultraluminous X-ray sources \citep[e.g.][]{Walton16}. % While WA/UFO signatures in the X-ray spectra are thought to be generic to accretion-powered sources, their launching mechanism is still poorly understood. % %Due to its highly ionized massive column, it seems less likely that acceleration mechanisms other than the action of magnetic fields could provide sufficient forces in general. On the other hand, in a situation where hard X-rays are somehow shielded to prevent gas from being overionized in the wind, the UV line opacity, possibly boosted also by soft X-rays, could be large enough to drive ions up to near-relativistic speeds as commonly thought in the case of UV BALs in distant QSOs (despite the known discrepancy that the estimated radiative power appears to be smaller than the observed kinetic power by orders of magnitude as discussed in \citealt{Matzeu16}). % \obj\ is an archetypal nearby ($z=0.184$), radio-quiet quasar (QSO) hosting a BH of mass $M \sim 10^9 \Msun$ \citep[e.g.][]{Reeves09}, being the most luminous AGN in the local universe with a bolometric luminosity of $L_{\rm bol} \sim 10^{47}$ erg/s. It is among the best studied QSOs for its strong UFO signatures observed in the Fe K band in the past X-ray observations with {\it Suzaku} \citep[e.g.][]{Reeves09,Reeves14,Gofford14,Matzeu16}, {\it XMM-Newton}/EPIC and {\it NuSTAR} \citep[e.g.][hereafter, N15]{Reeves03,Behar10,Nardini15}. A number of these extensive monitoring campaigns of \obj\ at different flux levels between 2001 and 2014 has provided us with insights into the variable nature of the UFOs %within individual observations and across different epochs since its first discovery \citep[e.g.][]{Reeves03}. Detailed spectral analyses have consistently confirmed that the highly ionized UFOs of $v/c \sim 0.2-0.3$ could be identified as the 1s-2p resonance transitions of \fexxvi\ with almost Compton-thick column of $N_H \lesssim 10^{24}$ cm$^{-2}$ and high ionization parameter of $\log \xi \sim 4-6$ \citep[e.g.][]{Reeves09}. % The distance of the observed UFOs is estimated to lie within $\sim 50 R_S$ where $R_S$ is the \sw radius %which fits within the expected UV/BLR emission radius %indicating a likely origin in the black hole (BH) vicinity \citep[e.g.][]{Reeves09}. The spectral variability has been studied as well suggesting that variable partial covering and/or intrinsic continuum variation may cause the observed short-term variability in X-ray \citep[][hereafter, M16]{Matzeu16}. \citet{Reeves18} also found an even faster Fe K UFO ($v/c \sim 0.42$) in 2017 {\it XMM-Newton/NuSTAR} observation. %this extremely luminous AGN has undergone a considerable spectral change that could presumably help suppress the otherwise present UFOs. The QSO outflows have been thought to be driven by radiation pressure by their O/UV flux, in a manner similar to that observed in massive stars \citep[][]{King10,Hagino15,KingPounds15,Hagino17,Nomura17}. However, most UFOs are so highly ionized that there is little UV or soft X-ray opacity, making this process very inefficient \citep[e.g.][]{Higginbottom14}. Alternatively, winds can be driven by the action of global magnetic fields threading the accretion disk to provide a plausible means of efficient acceleration as observed (e.g. \citealt{BP82}, hereafter, BP82; \citealt{CL94}, hereafter, CL94; \citealt{KK94}, \citealt{F10}, hereafter F10; \citealt{K12}; \citealt{K15}; \citealt{Kraemer18}). \obj\ UFOs are especially interesting because of the variation of its properties with the X-ray variability over the past fifteen or so years. For example, \citet[][hereafter, M17]{Matzeu17} noted, over a decade of multi-epoch observations, that the detected Fe K UFO velocity is well correlated with X-ray luminosity $L_X$ ($7-30$ keV) suggestive of a radiative-driven origin. Independently, \citet[][hereafter, Pa18]{Parker18} showed a likely anti-correlation between the observed UFO EW and $2-10$ keV flux. A similar trend is also reported in the {\it XMM-Newton}/EPIC-pn and {\it NuSTAR} spectra for the low-redshift narrow-line Seyfert 1 (NLS1), IRAS~13224-3809 (\citealt{Parker17}, hereafter P17; \citealt{Pinto18}, hereafter Pi18) possibly accreting at near-Eddington rate. Motivated by these findings, we investigate in this Letter the dependence of the variation of the UFO properties on the X-ray flux within the MHD wind framework of our earlier studies. In \S 2, we overview the proposed wind models. In \S 3, we present our results and demonstrate that magnetically-driven winds can account for the observed correlations. In \S 4 we conclude with a summary and discussion. %P17: nature, IRAS 13224-3809 %P18: PCA analysis showing EW correlations %M16; low-flux state showing intrinsic SED %M17: correlation with vout %\citep[][]{Peterson04}; it has been observed with {\it Chandra}/HETGS a number %of times to date since 2000 \citep[][]{Kaspi00}, and also in conjunction with %simultaneous observations by {\it ASCA} and {\it RXTE} \citep[][hereafter K01]{K01}. %% %It is one of the most intensively monitored Seyfert galaxies for its high-resolution %absorption study with a total duration of $900$ ks {\it Chandra} grating data %\citep[i.e. five $\sim 170$ ks observations and a $56$ ks one;][hereafter, K02]{K02}. %In addition to a series of ionized absorbers in X-ray, \obj\ is also known to exhibit %UV absorbers detected with {\it FUSE} and {\it HST}/STIS \citep[e.g.][and references %therein]{KraemerCrenshawGabel01}. In particular, K02 conducted an exploratory spectral %analysis of the detected X-ray absorbers identifying the physical characteristics of %individual ions in the broad-band stacked spectrum assuming phenomenological multiple %absorption systems. They did not find any correlation of their velocity shifts or their %FWHMs with ionization potentials. Along similar lines, \citet{Krongold03} %\citep[also][]{Netzer03} analyzed the same 900 ks {\it Chandra} grating spectrum with a %self-consistent photoionization model, assuming a simple geometry that consists of a %central source emitting an ionizing SED and clouds of gas intercepting our line of %sight. Employing {\tt cloudy} \citep[version 90.04;][]{Ferland13} to obtain the %individual clouds' ionization state, they were able to constrain the physical %parameters of various ions, i.e. their ionization parameter $\xi$, column $N_H$, %outflow velocity $v$ and internal turbulent velocity $v_{\rm turb}$. % %==== below old text % %Ionized outflows are a common feature of Active Galactic Nuclei (AGN), manifesting %themselves as blue-shifted absorption spectral lines. Approximately $50\%-70\%$ of all %Seyfert 1's exhibit such features in their UV spectra and a similar fraction in their %X-ray spectra \citep{CKG03}. The most likely process responsible for the outflowing %plasma ionization is photoionization by the AGN ionizing continuum; in this respect, %X-rays appear to be of broader utility in probing their properties since X-ray %transitions span a much larger range in photoionization parameter $\xi$ than the UV %ones. % %X-ray absorption features were discovered first in the {\em Einstein} QSO spectra %\citep[e.g.][]{Halpern84}, with more significant detections of K-shell absorption edges %due to \ovii\ (0.74 keV) and \oviii\ (0.87 keV) by {\it ROSAT} \citep{NandraPounds92, Nandra93, Fiore93, Turner93}. Later on, the %improved spectroscopic capabilities of {\it ASCA} confirmed the robust presence of %these edges in many luminous Seyfert AGNs \citep[e.g.][]{Reynolds97,George98}; %attributed to absorption by warm plasmas ($T \sim 10^6$ K), they have been thenceforth %referred to as Warm Absorbers (WA). With the much enhanced spectral resolution and %senstitivity of dispersive spectrometers onboard {\it Chandra} and {\it XMM-Newton}, it %became obvious that there exists a plethora of absorption lines of various %charge states of many elements; these span a wide range of ionization parameter $\xi %\equiv L_{\rm ion}/(n r^2)$ \citep[e.g.][]{CKG03,Blustin05,Steenbrugge05,McKernan07} where $L_{\rm ion}$ is the %ionizing (X-ray) luminosity, $n$ is electron number density at radius $r$. Their %columns lie in the range $10^{20} \lesssim N_H \lesssim 10^{22}$ cm$^{-2}$, their %ionization parameter between $-1 \lesssim \log \xi \lesssim 4$, their temperatures %between $10^4 < T <10^7$ and the exhibit line-of-sight (LoS) velocities of $v/c \lesssim 0.01$ %\citep[e.g.][]{ReynoldsFabian95}, implying distances $r \gsim 10^4$ \sw radii, %employing the Keplerian association between velocity and radius. % | We have employed the MHD accretion disk wind model introduced in our earlier works (F10, F17, F18) to account for the observed %velocity and EW correlations of the \fexxvi\ UFO in \obj. %The importance of this approach is %that it implements a theoretically specific model while observationally well tested. Our model does not include explicitly the effects of radiation pressure, as do models specifically built to consider radiatively-driven outflows \citep[e.g.][]{Higginbottom14,Hagino17,Nomura17}. These are expected to be significant in sources accreting close to their Eddington rate ($\dot m_a \simeq 1)$ like \obj. %We have attempted to mock-up their effects with the introduction of the parameter $s$ and the dependence of $\eta_w$ on $\dot m_a$. We have demonstrated by spectral analysis %of the 2013/2014 {\it XMM-Newton/NuSTAR} data that the observed \fexxvi\ UFO feature of \obj\ can be successfully reproduced within the framework the magnetically-driven disk-winds. Our MHD-wind model can also account for the observed correlations of the UFO velocity and EW with X-ray flux over multi-epoch data. The model assumes that the wind mass flux can increase faster than X-rays, a situation not unreasonable in such high mass-accretion object where radiation can be trapped and advected with the flow. {\bf Figure~\ref{fig:wind3}} shows the calculated streamlines along with wind density $n(r,\theta)$ and fiducial ionization parameter $\xi(r,\theta)$ in the poloidal plane. %The predicted UFO characteristics therefore is governed almost exclusively by photoionization and %magnetized wind structure. % %The model in this work is successful in reproducing the observed \fexxvi\ spectrum and the %correlations. % The near-Eddington luminosity of \obj\ is crucial in producing an increase in the UFO velocity with $f_{44}$ since the \fexxvi\ velocity is closely related to the local escape velocity. %from the radius at which the plasma with the proper ionization has been launched. An increase in this velocity implies that the ionization front responsible for the UFO moves radially inwards. This seems natural only when flows are close to Thomson-thick (a fact determined not by the X-rays but by the near-Eddington O/UV luminosity of \obj), {\bf which is conceivable as the O/UV photons are closely related to mass-accretion rate}. Of support of this notion are the observations of a near-Eddington narrow-line Seyfert AGN, IRAS~13224-3809, which exhibits similar UFO correlations \citep{Pinto18}. % {\bf To explore this, we have considered different density profiles and confirmed that the wind of $p=0.9$ indeed produces a smaller \fexxvi\ column (i.e. lower EW) than does the $p=1.5$ case as discussed in \S 3.2. } %To explore this idea further, we ran additional {\tt xstar} calculations with different %density profiles, i.e. steeper ($p=1.5$) and shallower ($p=0.9$) ones as a function of %distance $r$ (see in \S 2). Because X-rays are progressively reprocessed by the wind along %a LoS, the density slope actually makes a large impact on ionization structure. We find that the wind of $p=0.9$ indeed produces a smaller \fexxvi\ column (i.e. lower EW) than the $p=1.5$ case due to its larger plasma column that scatter more easily the incoming X-rays through the wind while lowering $\xi$, as discussed in \S 3.2. % Low/sub-Eddington AGNs to the contrary are not expected to have this effect, allowing the EW to increase with X-ray strength as depicted in {\bf Figure~\ref{fig:wind2}a}. {\bf Overall, we note that the exact slope of the calculated correlation is sensitive to the values of $s$ and $p$ as well as the X-ray SED, and the study of such dependences will be left as a future work. } %Despite the successful modeling in this work, the exact quantitative aspect of the %correlations is sensitive to the specific physical variables such as the injected SED, %its source geometry and wind structure, for example. In reality, the wind is probably %hybrid (i.e. radiative MHD) subject to (strong) radiation field especially in bright %AGNs such as \obj, which is deliberately ignored in this work \citep[see][]{Everett05}. %Our models, as an extreme case, indicate that even a magnetic force alone can explain %the multi-year variability of the observed UFOs in \obj. % %\textcolor{blue}{ As discussed in \S 3.2 the correlations of ({\bf %Fig.~\ref{fig:wind2}a}) rely on the fact that the increase in mass accretion fails to %produce more \fexxvi\ absorption due to slightly lower ionization and X-ray scattering, %which reduce its EW. The velocity accordingly increases as the production site shifts %radially inwards. %% %To explore further this idea, we run additional {\tt xstar} calculations with different %density profiles, i.e. steeper ($p=1.5$) and shallower ($p=0.9$) ones as a function of %distance $r$ (see in \S 2). Because X-rays are progressively absorbed by the wind along %a LoS, the density slope actually makes a large impact on X-ray scattering and %ionization structure. We find that the wind of $p=0.9$ indeed produces smaller \fexxvi\ %ionic column than the $p=1.5$ case due to its larger plasma column that scatter more %easily the incoming X-rays through the wind while lowering $\xi$, as discussed %in \S 3.2. Consequently, the EW is greatly reduced in that case as expected. %%despite the wind's larger (column) density available for synthesizing \fexxvi\ ion otherwise. %} % %\textcolor{blue}{As discussed in \S 2 the ionization parameter $\log \xi$ for $s=1$ %scaling only slowly decreases with $f_{44}$ as $\xi \propto f_{44}^{-1}$. Pi18 shows %that $\xi$ of the detected Fe K UFOs in IRAS~13224 slowly increases with X-ray %luminosity. However, such $\xi$ vs. $L_{\rm ion}$ correlation needs caution because %$\xi$ is not a direct measurement and it is heavily dependent on the input radiation %field which is much less constrained. } % %% %\textcolor{blue}{We note also that our calculated EW appears to saturate at lower %$f_{44}$ showing a plateau, which may be already seen in observations (see %M17,P17,Pi18), whereas the statistics is less reliable). In our model the predicted EW %is further reduced with decreasing $f_{44}$ as wind density also decreases, which %predicts a peak EW at a characteristic X-ray luminosity. Such a trend should be %searched for in more multi-epoch observations. } % % %an increase in the expected EW during lower X-ray phases (i.e. $5 \lesssim f_{44} \lesssim 10$) as shown in {\bf Figure~\ref{fig:wind2}a}. This occurs only at relatively low $f_{44}$ regime because the illuminating X-rays are in fact able to transmit through the outer part of the wind without being significantly scattered when mass-accretion rate $\dot{m}_a$ is not yet high enough to produce Compton-thick winds. This allows for the spectrum to be relatively broad extended towards lower velocity end as seen for $f_{44} = 6$ in {\bf Figure~\ref{fig:wind2}b}. Hence, more plasma at larger distances will be equally photoionized % %perhaps in the same way as the WAs are in nearby Seyferts, % %with increasing $f_{44}$, thus enhancing the EW accordingly in this lower $\dot{m}_a$ regime. In the analysis of \obj\ in P18, there appears to be a potential indication that the observed EW may indeed reach a maximum or saturated value at the low X-ray flux (i.e. lower $\dot{m}_a$) in the derived correlation, % %but the statistics is not sufficient to make a definitive conclusion. If this is indeed the case, then this trend at lower X-ray strength should naturally be expected in more moderately luminous Seyfert AGNs where $\dot{m}_a$ is not so high; e.g. NGC~4151, NGC~4051, Mrk~509, among others that are known to exhibit the UFOs. % %As the wind density becomes highly enough (i.e. $f_{44} \gtrsim 5$), on the other hand, the EW starts dropping due to scattering by the high column density of the interior part of the wind that prevents the wind from being sufficiently ionized to yield \fexxvi\ UFOs, as described in \S 3.2. % %More observations especially probing lower X-ray epochs would be very helpful to better test this prediction. %In relation to the Fe K UFOs in bright QSOs, \cite{Dadina18} has reported a robust detection of a similar UFO at $E \approx 9.2$ keV with $v/c \sim 0.3$ in an {\it XMM-Newton} observation from a radio-loud lensed quasar, MG~J0414+0534 ($z=2.64$). %%among the UFOs seen in the other lensed quasars at high-redshift \citep[e.g.][]{Hasinger02,Chartas03,Chartas07,Chartas09a,Chartas16}. %It is worth noting that the observed UFO EW in this QSO (EW $\sim 235$ eV) appears to be much narrower than that from typical Fe K UFOs being less likely to be an edge feature. %%Considering its estimated intrinsic luminosity of $L_{\rm 2-10~keV} \approx 3 \times 10^{44}$ erg~s$^{-1}$, %Setting aside the intrinsic characteristics of MG~J0414+0534, we speculate that the observed narrow UFO profile can be naturally accounted for by the current model when the wind is near-Compton-thick as demonstrated in {\bf Figure~\ref{fig:wind2}b} (see the one in red). %Interestingly, a 2017 {\it XMM-Newton/NuSTAR} observation of \obj\ seems to exhibit an even faster Fe K UFO component ($v/c \sim 0.46$) in addition to the {\it slower} one discussed here ($v/c \sim 0.2-0.3$) considered in this work when the source was in a low flux state \citep{Reeves18}. While a proper modeling of this new absorber is beyond the scope of this paper, we note here that there always exists high-velocity layers of the wind launched from the inner part of the disk in the context of the MHD-driven model. The present model can be further exploited to better understand these UFOs within a unique framework. %These authors have argued that the both features are physically more favored, among other possibilities, in view of separate \fexxvi\ Ly$\alpha$ absorbers. A proper modeling of this newly discovered UFO component is beyond the scope of this paper, but we would like to speculate here certain interpretations. Although it is quite plausible, as discussed in \cite{Reeves18}, that the slow and fast UFOs originate from different outflows, (1) Our photoionization calculations for 2013/2014 observations predict near-relativistic UFOs from \fexxvi\ Ly$\alpha$ and Ly$\beta$ lines at $E \sim 9$ keV and $E \sim 11$ keV, respectively, as is reported in \cite{Reeves18}. However, the calculated Ly$\beta$ line is typically weaker than the Ly$\alpha$ line due mostly to its weaker oscillator strength inconsistent with data. (2) It is also conceivable that the detected line with $v/c \sim 0.2-0.3$ might be actually the \fexxv\ line (He-like Fe) in the past and the 2017 observations while the faster component can indeed be attributed to \fexxvi\ (H-like) ion. In this case, the slower component from \fexxv, currently identified as \fexxvi\ line with $v/c \sim 0.2-0.3$, would exhibit faster velocity of $v/c \gtrsim 0.3$ with its column being roughly comparable to that of \fexxvi. %In this case, however, it would be puzzling why \fexxvi\ ion were not present in the previous high flux state while appearing in this low flux state. %Detailed state-of-the-art spectroscopic analyses of the X-ray UFOs will be made possible with the launch of %{\it XARM} in the coming years and later on by ESA's {\it Athena} through micro-calorimeter %observations. Fast X-ray outflows have also been detected in BH XRBs \citep[e.g.][]{Miller15,Miller16} that can be considered as ``scaled-down" AGN UFOs in this framework. With timescales much shorter than those in AGNs, fast XRB winds over many binary orbits with X-ray variability may provide us with another valuable clue to systematically understand those correlations as discussed here for \obj. The planned future missions, such as {\it XARM} and {\it Athena}, will be able to better constrain the enigmatic UFO properties with unprecedented energy resolution, perhaps leading to answering the ultimate question of its launching mechanism. | 18 | 8 | 1808.08935 |
1808 | 1808.03792_arXiv.txt | We analyse 8 years of \textsc{PASS} 8 \textit{Fermi}-LAT data, in the 60 MeV - 300 GeV energy range, from 30 \textcolor{black}{high Galactic latitude} globular clusters. Six of these globular clusters are detected with a TS > 25, \textcolor{black}{with NGC 6254 being detected as gamma-ray bright for the first time.} The most significant detection is of the well-known globular cluster 47 Tuc, and we produce a refined spectral fit for this object with a log parabola model. NGC 6093, NGC 6752 and NGC 6254 are fitted with hard, flat power law models, NGC 7078 is best fitted with a soft power law and NGC 6218 is best fitted with a hard, broken power law. This variety of spectral models suggests that there is a variety of $\gamma$-ray source types within globular clusters, \textcolor{black}{in addition to} the traditional millisecond pulsar interpretation. We identify a correspondence between diffuse X-ray emission in globular cluster cores and gamma-ray emission. This connection suggests that gamma-ray emission in globular clusters could \textcolor{black}{also} arise from unresolved X-ray sources or a relativistic electron population, perhaps generated by the millisecond pulsars. X-ray observations of further gamma-ray bright globular clusters would allow a functional relationship to be determined between diffuse X-ray and gamma-ray emission. | The 155 known (\cite{RN37}) globular clusters (GCs) are bound, spherical stellar systems. They are old (of the order of 10\textsuperscript{10} years), dust-free satellites of the Milky Way galaxy, characterised by dense cores of 100 to 1000 stars per cubic parsec and consequently high stellar encounter rates. GCs are noted for hosting low mass X-ray binaries and populations of millisecond pulsars (MSPs) which arise from binary interactions. MSPs are strong gamma-ray sources emitting gamma-rays through curvature radiation and electron / positron pair production cascades in their magnetospheres. GCs may also contain central intermediate mass black holes (IMBH; \cite{RN124}) or reside within dark matter (DM) halos (\cite{1984ApJ...277..470P}), and in this context the annihilation of DM in IMBH could produce gamma-rays (\cite{RN180}). It is thus not surprising that GCs can also be significant gamma-ray sources. The \textit{Fermi}-LAT has been surveying the entire gamma-ray sky in the energy range from 60 MeV to more than 300 GeV since its launch in 2008. \textit{Fermi}-LAT is a pair production instrument consisting of tracker, calorimeter and anti-coincidence detector modules. The tracker determines the direction of gamma-ray photons, the calorimeter measures the photon energy and the anti-coincidence detector vetoes false events caused by cosmic rays (\cite{RN195}). The LAT event level analysis framework has been refined since 2008 in successive data releases. \textsc{PASS} 6 was the initial data release, followed by \textsc{PASS} 7 in 2011, which improved the gamma-ray Galactic diffuse emission model, instrument response functions and direction reconstruction above 3 GeV. The latest data release is \textsc{PASS} 8, which is a complete reworking of the dataset which reduces gamma-ray background, increases instrument effective area, improves the point spread function and allows analysis down to 60 MeV. The first gamma-ray detected GC was 47 Tuc, with around 8 months of \textit{Fermi}-LAT (large area telescope) observations providing a detection with a significance of 17 $\sigma$ (\cite{RN196}). Later \textit{Fermi}-LAT observations show that 47 Tuc exhibits an exponential cutoff power law spectrum between 200 MeV and 10 GeV (\cite{RN196} and \cite{RN176}) while GC NGC 6093 (M80) has been identified as a possible gamma-ray source and a possible detection of NGC 6752 has been confirmed (\cite{RN50}). More recently NGC 6218 and NGC 7078 have been detected using \textit{Fermi}-LAT \textsc{PASS} 8 data (\cite{RN2}). However, the previous analyses of 47 Tuc, NGC 6093 and NGC 6752 were performed using just the 2 years of \textit{Fermi}-LAT \textsc{PASS} 6 data then available. Thus, the published SED for 47 Tuc is coarsely binned at 4 bins per decade of energy (\cite{RN176}), the possible detection of NGC 6093 has not been refined and no spectral energy distribution has been produced for NGC 6752 \cite{RN50}. To date 25 GCs are known gamma-ray emitters (\cite{RN56}) and a re-survey with the most up to date \textsc{PASS} 8 \textit{Fermi}-LAT data is likely to refine spectra further and possibly make fresh detections due to the 1-7 years of further photon statistics since the last publications and the improved effective area of the LAT instrument. In addition, the latest \textsc{PASS} 8 data release and tools of the \textit{Fermi}-LAT, now allow spectral analysis in the 60 - 100 MeV range. This paper presents a new analysis of 30 GCs, and is structured as follows. In Section \ref{sec:GCSelection} we describe the selection criteria which result in the identification of the GCs for analysis. In Section \ref{sec:Analysis} we describe our analysis method for the detection of GCs, and variability, extension and spectral analyses for the GCs we detect. In Section \ref{sec:Results} we provide spectral models, spectra, light curves and flux determinations for the detected GCs along with their test statistic (TS) maps. For undetected GCs we present photon and energy flux upper limits. In Section \ref{sec:Discussion} we discuss whether the detected GC gamma-ray emission can be accounted for by MSPs on spectral grounds and determine a relationship between gamma-ray luminosity and diffuse X-ray luminosity in GCs. Finally in Section \ref{sec:Conclusion} we summarise our findings and make suggestions for future work. | \label{sec:Conclusion} We analyse 8 years of \textsc{PASS} 8 \textit{Fermi}-LAT data from 30 globular clusters. We refine the gamma-ray spectra of 5 previously detected GCs and detect NGC 6254 for the first time. NGC 6752 lacks detectable gamma-ray emission above 4 GeV, suggesting that this is the one object with emission dominated by MSPs. However, the spectral shapes of NGC 6093, NGC 6254, 47 Tuc and NGC 6218 suggest that other sources apart from MSPs contribute to the gamma-ray emission of these GCs. We also note that variability of NGC 6218 in particular points to a contribution other than MSPs to the gamma-ray emission, possibly a background AGN. An attempt to correlate the gamma-ray flux with the number of known MSPs in the GCs in our study reinforces this view, though there is some evidence for a correlation between the gamma-ray flux and the mass $\times$ normalised encounter rates. We note the presence of a link between GC core diffuse X-ray emission and GC gamma-ray emission. The core diffuse X-ray emission could be due to either unresolved point sources or to relativistic electrons in the GCs. In the latter case, one might expect TeV emission from the GCs due to inverse Compton radiation, and observations with ground-based gamma-ray telescopes such as CTA could resolve this issue. The link between core diffuse X-ray emission and gamma-ray emission is tentative, largely because there are relatively few gamma-ray emitting GCs for which X-ray observations are available. Further X-ray observations of GCs would be helpful in this regard. \begin{figure*} \includegraphics[width=\textwidth]{REV_MSP_VS_ER} \caption{ \textcolor{black}{Plot of inferred MSP count vs Encounter Rate $\Gamma$\textsubscript{e} from \citeauthor{RN101} which we renormalise assuming 47 Tuc encounter rate is 65 to allow comparison with \citeauthor{RN176}. Detected GC are labeled (red captions) and we show upper limits (black symbols) only for those GC at a helio-centric distance of less than 10.4 kpc which is the furthest detection in our study (NGC 7078). The line of best fit (ignoring upper limit values) is shown in green and has the functional form N\textsubscript{MSP}=(0.66$\pm$0.03)$\Gamma$\textsubscript{e} -(4.99$\pm$0.10)} } \label{fig:REV_MSP_VS_ER} \end{figure*} \begin{figure*} \includegraphics[width=\textwidth]{REV_METALLICITY} \caption{ \textcolor{black}{Plot of log gamma-ray luminosity vs Metallicity [Fe/H] for GCs with heliocentric distance of 10.4 kpc or less. Detected GCs are captioned in red whilst undetected GCs are captioned in black and shown as upper limits. There is no apparent correlation between luminosity and metallicity for the detected GCs in our study and NGC 6093 and NGC 7078 are the clear outliers. }} \label{fig:REV_METALLICITY} \end{figure*} \begin{figure*} \includegraphics[width=\textwidth]{MSP_COUNT} \caption{ Log Plot of gamma-ray luminosity vs number of known MSPs in each GC. We show GCs which are detected in this study (in red) and undetected GCs (in black) with known MSPs. For undetected GCs an upper luminosity limit is shown} \label{fig:MSP_COUNT} \end{figure*} \begin{figure*} \includegraphics[width=\textwidth]{ENCOUNTER_RATE_LUMINOSITY} \caption{ Log plot of gamma-ray luminosity vs product of cluster mass and a normalised encounter rate from \citeauthor{RN101} for GCs detected in this study (in red) and selected GCs from \citeauthor{RN56} (in black) with known masses and encounter rates. In \citeauthor{RN101} the encounter rate of 47 Tuc is set to an arbitrary value of 1000 and the encounter rate of other GCs are determined relative to that of 47 Tuc. The line of best fit (ignoring upper limit values) is shown in green and has functional form Log (L\textsubscript{$\gamma$})=0.30 Log (Mass*Encounter Rate) + 33.7 } \label{fig:ENCOUNTER_RATE_LUMINOSITY} \end{figure*} \begin{figure} \includegraphics[width=\textwidth]{Lightcurves/NON_VARIABLE} \caption{Lightcurves for the 6 detected GCs. Only NGC 6218 has some evidence of variability on 6 month timescales at the 3 $\sigma$ level } \label{fig:NON_VARIABLE} \end{figure} | 18 | 8 | 1808.03792 |
1808 | 1808.04384_arXiv.txt | {The process of density field reconstruction enhances the statistical power of distance scale measurements using baryon acoustic oscillations (BAO). During this process a fiducial cosmology is assumed in order to convert sky coordinates and redshifts into distances; fiducial bias and redshift-space-distortion parameters are also assumed in this procedure. We analytically assess the impact of incorrect cosmology and bias assumptions on the post-reconstruction power spectra using low-order Lagrangian perturbation theory, deriving general expressions for the incorrectly reconstructed spectra. We find that the BAO peak location appears to shift only by a negligible amount due to wrong assumptions made during reconstruction. However, the shape of the BAO peak and the quadrupole amplitude can be affected by such errors (at the percent- and five-percent-level respectively), which potentially could cause small biases in parameter inference for future surveys; we outline solutions to such complications.} \arxivnumber{18MM.NNNNN} \begin{document} | The measurement of baryon acoustic oscillations in the clustering of galaxies and the intergalactic medium provides some of our tightest constraints on the cosmological distance scale (for a review of recent measurements see \cite{BOSS_DR12}, e.g.~Fig.~14, for reviews of the methods see \cite{Wei13,PDG14}). With the exception of the intergalactic medium results, all recent measurements have employed a process known as ``reconstruction'' \cite{ESSS07}, which aims to undo some of the loss of signal due to non-linear evolution by estimating the motion of tracers under gravity and reversing it. In order to perform reconstruction, one typically assumes a fiducial cosmology to convert sky positions and redshifts into comoving coordinates; one also assumes a fiducial bias and redshift-space distortion (RSD) parameter to derive the density field from the galaxy distribution. However, the cosmology assumed may not match the true, underlying cosmology, and the fiducial bias and RSD parameters may also not equal their true values. In this paper we use Lagrangian perturbation theory to assess the impact of such errors in the assumed cosmology, biasing and RSD upon the two-point clustering of the reconstructed field. Reconstruction has been extensively studied in the literature using both N-body simulations and analytic models \cite{ESSS07,PWC09,NWP09,Seo10,Pad12,TasZal12b,McCSza12,SheZal12, Xu13,Sch15,BPH15,AchBla15,Whi15a,Coh16,Seo16,Var16,Hik17}. While the impact of incorrect choices of distance, growth rate and bias has been studied in numerical simulations in particular cases, we are not aware of any analytic treatment. Since the standard algorithm is based upon the Zel'dovich approximation \cite{Zel70}, we choose to similarly use low-order Lagrangian perturbation theory to assess the impact of incorrect assumptions about the cosmological model upon the recovered statistics. While the results will only be approximate, due to our approximate model of structure formation, they allow us to provide general formulae and hence estimate the errors introduced in any scenario. The outline of the paper is as follows. In Section \ref{sec:background} we briefly review Lagrangian perturbation theory and reconstruction and define our notation. Section \ref{sec:distances_wrong} explains the calculation of the reconstructed power spectra in a simple case where only the assumed distances are incorrect. Section \ref{sec:everything_wrong} describes the more general calculation where the distances, the growth rate and the bias assumed can all be wrong. In Section \ref{sec:discussion}, we evaluate our calculations and discuss the results. We conclude in Section \ref{sec:conclusions}. | \label{sec:conclusions} In this paper we have analytically calculated how wrong assumptions about the fiducial cosmology, bias and RSD parameters impact the reconstructed power spectrum. We have derived a full expression for this power spectrum (Eq.~\ref{eqn:Pfalse_rec}) and have evaluated it for realistic scenarios. Assuming large but not unrealistic reconstruction errors ($3\%$ on distances, $10\%$ on bias, $30\%$ on RSD) we find that: \begin{itemize} \item The shifts in the BAO monopole peak position(s) in both Fourier and configuration space appear negligible. \item The shape of the BAO peak or the oscillation envelope in Fourier space can be modified at the percent level. Fitting with a wrong template could potentially lead to small errors in parameter inference; possible solutions include incorporating a wrong-reconstruction template (from Eq.~\ref{eqn:Pfalse_rec}) in the likelihood, marginalizing over an oscillation envelope shape in a Fourier space analysis, or potentially iterating the reconstruction. \item The quadrupole can be affected at the $5\%$ level by even moderate errors in the assumed cosmology, bias and RSD. \end{itemize} Future research directions include investigating new likelihoods or iterative fitting and examining whether wrong reconstruction impacts BAO constraints on $N_\mathrm{eff}$, the effective number of neutrinos \cite{NeffTh,NeffMeas}. | 18 | 8 | 1808.04384 |
1808 | 1808.00044_arXiv.txt | { Formulation of the Lagrangian approach is presented for studying features of motions of stellar bodies with non-zero rest mass in the vicinity of fast-spinning black holes. The structure of Lagrangian is discussed. The general method is applied to description of body motion in the Kerr model of space--time to transition to the problem of tidal disruption of elastic bodies by strong gravitational field. } \keyword{Fast--spinning black hole, Kerr space--time, Structure of Lagrangian, Motion regimes } \PACS{04.25.-g; 98.10.+z} \begin{document} | \label{se1} In contrast to the well-known ray-method for describing (in non-flat space-time) motion of both photon and "test body" -- an idealized conceptualization of a material object with non-zero but small mass (not perturbing space-time around it) -- we formulate a step-by-step Lagrangian approach that not only describes the motion of the test body with non-zero rest mass (obviously, it must move along a geodesic), but also describes interactions of test--like bodies in a multi-body system, for example, a system of interacting bodies moving in a given strong gravitational field. \begin{figure}[h!] \centering \includegraphics[width=8cm]{galactic-center-ucla-2014.eps} % \caption{ There exist indications that the central zone of our galaxy contains a compact, possibly rotating, object with mass $\sim 4 \times 10^6 M_\odot$. This super--massive object, thought to be a black hole, exerts strong influence on the dynamics of nearby stars. (See, for example, \citet{g14, r08}, as well as \citet{kt14}.) The image (from \citep{g14}) shows tracks of the brightest stars near the center of the Milky Way. The image presented here was obtained by UCLA research team based on data obtained by the teams at the W.M.Keck Observatory over the span of two decades. The orbits plot as dots the star positions at one-year intervals. The central arrow (added by us) points at the presumed location of the super--massive black hole. } \label{bhMW} \end{figure} The need for the ability to conduct detailed examinations of such scenario is obvious. One vivid example of its importance comes from the astronomical observations showing that at the center of our galaxy a super--massive black hole significantly affects the dynamics of nearby stars (see Fig.\ref{bhMW}). This black hole is apparently one of the closest such objects to us. It represents an extraordinary, natural, laboratory not only for validating the general relativity theory overall, but also for testing specific theoretical models by comparing them with direct observational data. Further improvements in observational techniques combined with continuous (rather than periodic or ad hoc) collection of data can generate invaluable material for fundamental research. In view of this, the importance of the task of formulating the proper procedure for describing motions of bodies, both non-deformable and such that can be torn apart by tidal forces, is apparent. The framework of the Lagrangian Approach with the relativistically invariant Lagrangian is best-suited for the task. The structure of the articles is as follows. In Sec.~2, we lay out the essentials for working with the space-time metrics near a \emph{rotating} black hole. This section may be skipped by an experienced reader, but it is essential for the completeness of the presentation. In Sec.~3, we recall the essence of the principle of least action and set up the model of the relativistically invariant Lagrangian for describing motions of bodies near fast-rotating black holes. Sec.~4 discusses the Euler-Lagrange motion equations for bodies with non-zero mass, that are derived from this principle. The special case of planar parabolic motion of bodies is considered in Sec.~5. In Sec.~6, we discuss the obtained expression for Lagrangian in the post-newtonian approximation, which serves as the (complexity-reducing) basis for numerical simulations of tidal disruptions of elastic, plasma, "droplet-like", and so on, bodies, when they enter the close vicinity of a fast-rotating black hole. Sec.~7 concludes with discussion. | This paper was envisioned with dual goals in mind: to obtain concrete results and to highlight some methodological aspects involved in the finding of these results. The material concerning the space-time metric near black holes, coordinate choice, etc., contained in the introductory part of the paper, can be found in works on general relativity, at least implicitly. However, these foundation elements are scattered (sometimes in implicit form) through various works (cited in our bibliography and in references therein). We aggregated this dispersed material. An important purpose of this paper is to illustrate the "Lagrangian" method involved in the conducted study. We wished to convey to the readers who desire fresh perspectives and are unencumbered by ingrained authoritative dogmas, the idea that Method is more important than Result.\footnote{ L. Landau: "A method is more important than a discovery, since the right method will lead to new and even more important discoveries." \url{http://www.azquotes.com/quote/1263664}} In this paper, we laid out a step-by-step approach, based on the relativistically invariant, "square--root--form", Lagrangian Eq.~(\ref{i3d}), for studying problems of compact and compound stellar body motion in the vicinity of rotating black holes. We presented a classification of possible regimes of motion, and demonstrated how a body, once in the vicinity of a rotating black hole, becomes involved in the whirlpool--like rotation of the space-time (Fig. \ref{zakhvat}). We determined the conditions for the body's capture and involvement in the co-rotative motion with the black hole. We showed that in describing interaction between components of a multi-body system moving near the black hole, one can use Lagrangian in the form of Eqs.~(\ref{LagFin})--(\ref{Lagn}) derived in the post-Newtonian approximation with respect to parameter $r_g / r$. Obviously, bodies traveling along curved trajectories must emit gravitational waves (\cite{ll69}, \S 110) and lose energy in the process. This energy loss per unit of time is noticeable only at the fifth order of magnitude in the Lagrangian expansion with respect to small parameter $\sim (v/c)$. With respect to the terms with the first four orders, the energy of the body remains constant. This means that in the expression for the Lagrangian, it suffices to include the terms up to the fourth degree of smallness with respect to $c^{-1}$, and not be burdened with more precise calculations. For a metric tensor written in the form $g_{\alpha \beta} = \eta_{\alpha \beta} + h_{\alpha \beta}$ (here $\eta_{\alpha \beta}$ is Minkowski tensor), Lagrangian $L = - m c ((\eta_{\alpha \beta} + h_{\alpha \beta}) u^\alpha u^\beta )^{1/2}$ with $u^\gamma = (1, \dot{q}^i)$. Here, the affine parameter $\tau$ is the time measured by the remote observer, i.e. $\tau = t$, and dot denotes the derivative with respect to this time. Obviously, $\eta_{\alpha \beta} u^\alpha u^\beta = c^2 - v_j v^j$ with $v_j v^j = \textbf{v}^2$, where $\textbf{v}$ is the usual 3-velocity in the 4-coordinate system of the remote observer. Therefore, the Lagrangian for one body in the (given) gravitational field takes form \begin{eqnarray} \label{fn1} L = - m c \frac{d s}{d t} = - m c^2 \bigg( 1 + h_{00} + 2 h_{0 i} \frac{v^i}{c} - \frac{v^2}{c^2} + h_{ij} \frac{v^i v^j}{c^2} \bigg)^{1/2} \, . \end{eqnarray} Expanding the radical into series up to and including $c^{-4}$, yields: \begin{eqnarray} \label{fn2} L = - m c^2 + \frac{m v^2}{2} + \frac{m v^4}{8 c^2} - m c^2 \bigg( \frac{h_{00}}{2} + h_{0 i} \frac{v^i}{c} + \frac{1}{2} h_{i j} \frac{v^i v^j}{c^2} - \frac{h_{00}^2}{8} + \frac{h_{00}}{4} \frac{v^2_a}{c^2} \bigg) \, . \end{eqnarray} This explains the choice of approximation~(\ref{i6g}). In the context of the presented framework and findings, a question may be asked: what additional information can be extracted from the observational results presented in Fig.~\ref{bhMW}, beyond the conclusion that bright stellar objects orbit around a massive compact object with $M \sim 4.1 \times 10^6 M_\odot$, presumably a black hole? Let's make a few quantitative estimates. The Schwarzschild radius, $r_g = 2 G M / c^2$, for a body with mass $M$ is of order $\simeq 2.95 (M / M_\odot) \, km $, i.e., for the black hole in our galaxy, $r_{BH} \simeq 10^7 \, km \simeq 0.07 \, AU$. The moving objects in Fig.~\ref{bhMW} follow approximately elliptic trajectories. Only two trajectories are closed, i.e., objects SO-2 and SO-102 have completed their entire rotations over the duration of monitoring. The semi-major axis, $a$, of an elliptic orbit and the distance of the closest approach (periapse), $r_{min}$, with the gravitating center $M$, are related as $a = r_g / u_1 (1 - \epsilon)$, where $\epsilon$ is orbit eccentricity and $u_1 = r_g / r_{min}$. For a star that does not approach the black hole too closely, simple estimates, following from the classical celestial mechanics, can be used instead of cumbersome formulae of the relativistic theory.\footnote{ In general, however, the possibility that "dark matter" may have an impact on the motion of stars in the central zone of our Galaxy should also be noted and considered (see, for example, \citet{tp12} and references therein). } Then, from the law $a^3 / T^2 = (8 \pi^2)^{-1} r_g c^2$, follows the relationship between characteristic time/period and the distance of the closest approach -- $T^2 (1 - \epsilon)^3 = (8 \pi^2) u_1^{-3}(r_g / c)^2$. Numerically, $T = 9.461 \times 10^{12} \tau $, where $T$ is measured in seconds and $\tau$ in years. Then we obtain relationship $\tau^2 (1 - \epsilon)^3 u_1^3 = 7.72 \times 10^{-24} (M_h / M_\odot)^2$, which obviously is strictly valid only for the newtonian mechanics, but can be also used for estimations even in post--newtonian models. This relationship permits not only to estimate $M_h$ based on measured $\tau, \epsilon, u_1$, but also to assess (in terms of the order of magnitude) which values of $\epsilon$ are acceptable for experimentalists, so that the effects (similar to those shown in Fig.~\ref{ProlTraj}) of the general relativity theory become apparent. For regimes with parameter $u_1 \geq 0.2$ for the observed stars during $\tau \sim 0.1 \div 1$ (about a few months or a year), the above-mentioned relationship indicates that the eccentricity of the analyzed orbits must be not small. To the contrary, parameter $\epsilon$ must be close to $1$, $\epsilon \rightarrow 1$, i.e. trajectories of the observed stars should resemble either very elongated ellipses, or fly-through, quasi-parabolic paths. The movement of stars near the black hole would then occur very quickly, even taking into account the red-shift affecting the distant observer. Therefore, high-resolution observations of "bright" moving objects in the region shown in Fig.~\ref{bhMW} must be conducted continuously, not occasionally or with long pauses. At present, several teams conduct observations of the motion of stars orbiting around the supermassive black hole in our Galactic center that permit probing the gravitational theory (\citet{g14}, \citet{peswhpgdfogg17}, \citet{hd17}, \citet{peskzzs17}, \citet{zly15}). In view of their estimations (see for example \citet{peskzzs17}) the actually measurable value of parameter $u_1$ (the inverse periapse normalized by the Schwarzschild radius) is of order $7 \times 10^{-4}$ (for object S2 in the S-star cluster in the Galactic center region, which has the shortest period). Therefore, observation techniques need to evolve further to be able to capture stellar motions with parameters close to the above-mentioned $u_1$ and $\tau$. The challenges of obtaining reliable experimental data, which may escaped attention of theoreticians and numerical modelers, can be clearly seen in Fig.~\ref{S2andS38} (from\citet{bgsmyambdlmmsw16}). \begin{figure}[h!] \centering \includegraphics[width=9cm]{orbits.eps} \caption{Positions on the plane of the sky and best-fit orbits of objects S0-2 (blue) and for S0-38 (red) for the period from 1995 to 2014. From\citet{bgsmyambdlmmsw16}. } \label{S2andS38} \end{figure} Fig.~\ref{S2andS38} shows the positions of the two stellar objects, S0-2 (blue) and for S0-38 (red), from 1995 to 2014. The lines show best-fit orbits for the stars. Both stars orbit clockwise on the plane of the sky. For S0-38 (red), some observations show greater uncertainty ranges due to the star's proximity to other objects. In general, equations following from Lagrangians~(\ref{LagFin}) and (\ref{Lagn}) can be useful for studying the black hole's tidal effects onto elastic bodies and analyzing the bodies' potential destruction, as illustrated in Fig.~\ref{TidalDisrupt}. Furthermore, future discoveries of "bright" objects traveling along \emph{fly-through} trajectories close to this super-massive black hole (similar to those shown in Fig.~\ref{ProlTraj}) would serve as convincing \emph{direct} evidence for the general relativity theory and as an additional confirmation of the existence of black holes. \vspace{6pt} \conflictsofinterest{The authors declare that there is no conflict of interests regarding the publication of this article.} \reftitle{References} | 18 | 8 | 1808.00044 |
1808 | 1808.05985_arXiv.txt | To complement the 2MASS Redshift Survey (2MRS) and the 2MASS Tully-Fisher survey (2MTF) a search for 21cm \HI\ line emission of 2MASS bright galaxy candidates has been pursued along the dust-obscured plane of the Milky Way with the 100m \nan\ Radio Telescope. For our sample selection we adopted an isophotal extinction-corrected \K -band magnitude limit of $K_{\rm s}^{\rm o} = 11\fm25$, corresponding to the first 2MRS data release and 2MTF, for which the 2MASX completeness level remains fairly constant deep into the Zone of Avoidance (ZoA). About one thousand galaxies without prior redshift measurement accessible from \nan\ (Dec $> -40\degr$) were observed to an rms noise level of $\sim 3$\,mJy for the velocity range $-250$ to $10\,600$\,\kms. This resulted in 220 clear and 12 marginal detections of the target sample. Only few detections have redshifts above 8000\,\kms\ due to recurring radio frequency interference (RFI). A further 29 detections and 6 marginals have their origin in non-target galaxies in the telescope beam. The newly detected galaxies are on average considerably more \HI-rich (mostly $10^9 - 10^{10}$M$_\odot$) compared to systematic (blind) \HI\ surveys. The \HI\ detections reveal various new filaments crossing the mostly uncharted northern ZoA (e.g. at $\ell \sim 90\degr, 130\degr, 160\degr$), whilst consolidating galaxy agglomerations in Monoceros and Puppis ($\ell \sim 220\degr, 240\degr$). Considerably new insight has been gained about the extent of the Perseus-Pisces Supercluster through the confirmation of a ridge ($\ell \sim 160\degr$) encompassing the 3C129 cluster that links Perseus-Pisces to Lynx, and the continuation of the second Perseus-Pisces arm ($\ell \sim 90\degr$) across the ZoA. | 18 | 8 | 1808.05985 |
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1808 | 1808.02041_arXiv.txt | We present initial results from a deep neutral hydrogen ($\hi$) survey of the HALOGAS galaxy sample, which includes the spiral galaxies NGC891, NGC925, NGC4414, and NGC4565, performed with the Robert C. Byrd Green Bank Telescope (GBT). The resulting observations cover at least four deg$^2$ around these galaxies with an average 5$\sigma$ detection limit of 1.2$\times$10$^{18}$ cm$^{-2}$ over a velocity range of 20 km s$^{-1}$ and angular scale of 9.1$'$. In addition to detecting the same total flux as the GBT data, the spatial distribution of the GBT and original Westerbork Synthesis Radio Telescope (WSRT) data match well at equal spatial resolutions. The $\hi$ mass fraction below $\hi$ column densities of 10$^{19}$ cm$^{-2}$ is, on average, 2\%. We discuss the possible origins of low column density $\hi$ of nearby spiral galaxies. The absence of a considerable amount of newly detected $\hi$ by the GBT indicates these galaxies do not have significant extended diffuse $\hi$ structures, and suggests future surveys planned with the SKA and its precursors must go \textit{at least} as deep as 10$^{17}$ cm$^{-2}$ in column density to significantly increase the probability of detecting $\hi$ associated with the cosmic web and/or cold mode accretion. | \setcounter{footnote}{0} Resolved neutral hydrogen ($\hi$) observations undertaken over the past decade have revealed many intricate details related to the morphology and dynamics of spiral galaxies. A primary science goal of recent large surveys is to develop a deep understanding of how physical processes within the disks of spiral galaxies, such as star formation and the subsequent stellar feedback, affect their local circumgalactic environments. Surveys such as The $\hi$ Nearby Galaxy Survey (THINGS; \citealt{walter08}) and Hydrogen Accretion in LOcal GAlaxies Survey (HALOGAS; \citealt{heald11}; hereby referred to as H11) performed with the Very Large Array (VLA) and Westerbork Synthesis Radio Telescope (WSRT), respectively, provide high resolution maps of the environments around nearby spiral galaxies. Accretion of diffuse gas onto the disks of galaxies from the intergalactic medium (IGM) is a possible explanation for how the $\hi$ content of galaxies has remained relatively constant since $z\sim$ 2 while the star formation rate was up to 10 times higher at high redshifts \citep{not12,madDickinson14}. The constant $\hi$ content implies that galaxies have somehow replenished themselves with enough gas to fuel continuous star formation. And though not directly responsible for star formation, $\hi$ is an intermediate phase towards molecular hydrogen, which is the raw ingredient of the star formation fuel. If the star formation is to continue, external gas has to be accreted and pass through the $\hi$ phase at some stage in the accretion process. Observationally inferred accretion rates as traced by $\hi$, however, fall between 0.1 and 0.2 $\Msun$ at low redshifts. This is a full order of magnitude lower than what is needed for galaxies to continually form stars at their current rates \citep{sancisi08,kauff10}. This discrepancy presents two intriguing scenarios: the cycle of star formation will eventually exhaust all of the available fuel within a few Gyr and star formation itself will gradually cease, or processes that refuel galaxies with the necessary gas have been missed by previous surveys. Numerical simulations have shown a likely mechanism for refueling star formation is through a quasi-spherical `hot' mode and filamentary `cold' mode \citep{keres05,keres09,birnDek03}. Cold in the context of these numerical simulations refers to gas that has not been heated above the virial temperature of the galaxy's potential well ($\sim$ 10$^5$ K), and hot refers to gas that has virialized in a process akin to the classical theory of galaxy formation in which shock-heated, virialized gas with short cooling timescales accretes onto the central galaxy (e.g., \citealt{reesOs77}). These simulations also suggest cold mode accretion was the dominant form of accretion at $z \ge$ 1 for all systems, and remains prevalent through $z=0$ for galaxies in low-density environments ($n_{gal}$ $\lesssim$ 1 h$^3$ Mpc$^{-3}$) and $M_{halo}$ $\lesssim$ 10$^{11.4}$ $\Msun$ (or $M_{bary}$ $\leq$ 10$^{10.3}$ $\Msun$). For perspective, our own Milky Way has a virial (and thus halo) mass on the order of 10$^{12}$ $\Msun$. These cold flows should exist in the form of vast filaments of cold, diffuse gas that permeate through the hot halo \citep{keres05}. Comparisons by \citet{nelson13} between the smoothed particle hydrodynamic (SPH) numerical scheme employed in \citet{keres05,keres09} and more sophisticated adaptive mesh refinement (AMR) simulations revealed the relative contribution of the cold mode is likely overestimated in earlier SPH simulations due to inherent numerical deficiencies. Nevertheless, the AMR simulations do show $\textit{some}$ fraction of gas is accreted cold. The temperature of the gas in these predicted cold filaments is too high for a significant amount of neutral gas to exist within the largely ionized medium. However, AMR hydrodynamic simulations presented by \citet{joung12} show large amplitude non-linear perturbations can create cooling instabilities in which gas is collisionally excited and cools through subsequent radiative de-excitation of excited states. Large filamentary flows of inflowing gas are a possible seeding mechanism for non-linear perturbations, which allow gas to cool enough to form $\hi$ clouds within the inner most regions of the halo (R $\leq$ 100 kpc) at $\hi$ column densities ($N_{HI}$ $\leq$ 10$^{18}$ cm$^{-2}$) currently detectable with existing telescopes. More recent independent ballistic models show galactic fountain activity can account for the presence of extraplanar $\hi$ around the Milky Way (in clouds like Complex C; \citealt{frat15}) and NGC891 \citep{frat06}. In addition, \citet{frat17} describes the condensation of hot coronal gas in the wake of the interaction with cooler galactic fountain gas, showing that fountain driven accretion can cool enough lower coronal gas to sufficiently extend the gas depletion time. Observational evidence for predicted cold flows is very limited. Though absorption measurements, \citet{stocke10} and \citet{ribaudo11} both find low metallicity gas infalling onto a nearby solar metallicity Lymann Limit System whose mass is consistent with the presence of cold flows predicted by simulations. The presence of infalling, low metallicity gas is certainly consistent with cold flows, but these measurements do not reveal any information about the extended spatial distribution of the accretion. Absorption measurements are very promising in that they accurately probe the metallicity of galaxy halos, but such studies require a quasar or other bright background source to measure the absorption line of interest. Absorption studies of the Milky Way, in which these desired sightlines are abundant, show our own Galaxy is surrounded by an immense amount of low column density gas that is both ionized and neutral (e.g., \citealt{wakker03,richt17}) with temperatures ranging from 10$^2$ to 10$^7$ K. Detection in emission does not rely on serendipitous sightlines required for external galaxies, and will constrain the large scale extent of the predicted cold flows or a potential diffuse component. The unrivaled point source response of radio interferometers like the WSRT and VLA allows for incredible high resolution mapping capabilities at angular resolutions $\sim$$\frac{\lambda}{b_{max}}$, where $b_{max}$ is the maximum baseline ($b_{max} = 2.7$ km for the WSRT), which reveals the small scale structure of galaxies. On the downside, interferometers act as spatial filters by construction, and in particular due to the minimum possible spacing between neighboring telescopes in an interferometer (i.e., the physical size of each dish), there is a gap in $\it{u-v}$ coverage at large angular scales from the absence of short baselines. This gap is often referred to as the `short-spacing' problem, and it limits the amount of large scale structure an interferometer is able to detect (e.g., \citealt{brWl85}). As a consequence of the lack of sensitivity at large angular scales, past $\hi$ observations performed with interferometers may have missed significant reservoirs of gas around galaxies. On the other hand, the $\hi$ is observed in channels covering small velocity ranges only, and thus potentially does not extend enough to cause the sampled baselines to miss several interesting low-density, diffuse features. The full $\it{u-v}$ coverage capability of single dish telescopes \citep{stan02} permits the detection of structure at all angular scales to test that notion. The unblocked aperture design of the GBT and resulting low sidelobes coupled with the compromise between resolution (9.1$'$) and high surface brightness sensitivity (T$_{sys}$ $\lesssim$ 20 K) make it the ideal instrument to look for low column density structure around the HALOGAS sources. The few surveys that have mapped down to $N_{HI}$ $\lesssim$ 10$^{19}$ cm$^{-2}$ have uncovered several interesting low density, diffuse features. Perhaps most notably, \citet{brth04} discovered a low column density $\hi$ filament connecting M31 and M33. Two possible explanations for its origin have been presented since its discovery: either it is similar to filaments seen in simulations of the cosmic web \citep{pop09}, and thus an observational example of the cold mode accretion process, or it was created via a past tidal interaction between M31 and M33 \citep{bek08, putman09}. Higher resolution observations with the GBT by \citet{wolfe13} and \citet{wolfe16} show that this filament is clumpy in nature and made up of small $\hi$ clouds with M$_{HI}$ $\sim$ 10$^{4-5}$ $\Msun$, $N_{HI}$ $\sim$ 10$^{18}$ cm$^{-2}$, and diameters on the order of kpc. M31 has a M$_{dyn}$ $\sim$ 1.3 $\times$10$^{12}$ $\Msun$ \citep{cor10}, which suggests the cold mode accretion scenario is unlikely. Furthermore, the total $\hi$ mass of these clouds is only 4.6$\times$10$^{6}$ $\Msun$ providing only meager neutral mass accretion rates for a conservative infall time estimates of 10$^{7-8}$ years. The origin of these clouds is still an open and intriguing question which can be answered by utilizing sensitive observations of the $\hi$ within the circumgalactic environment of M31 and M33 \citep{wolfe16}. Other recent detections by the GBT of large $\hi$ structures in NGC6946 by \citet{pisano14} and NGC2403 by \citet{deBlok14} suggest these features are seen around a variety of galaxies. In order to determine the true origin of these filaments, resolve the discrepancy between observed accretion rates and SFRs, and obtain a comprehensive understanding of how the disks of galaxies interact with their surrounding circumgalactic environment, a comprehensive $\hi$ census spanning a wide range of astrophysically interesting properties (e.g., dynamical mass, total $\hi$ mass, halo mass, SFRs, etc) is required. A complete census of these properties will build up large number statistics and uncover any underlying correlations between intrinsic galaxy properties and possible signatures of accretion. The HALOGAS observations of 24 nearby galaxies obtained with the WSRT and the THINGS survey with the VLA are critical steps towards just such a census. To ensure this census is absolutely complete, interferometer observations must be supplemented with large single dish observations to cover all angular scales to assure large-scale emission is not resolved out by interferometers, and to map down to the lowest possible column density levels. In this pilot paper we present data and analysis from four sources from the HALOGAS survey: NGC891, NGC925, NGC4414, and NGC4565. These GBT maps are among the deepest ($N_{HI}$$\sim$10$^{18}$ cm$^{-2}$) for external galaxies obtained to date in $\hi$. This paper serves as an introduction to the full survey as a way to outline our analysis methods and highlight our overall goals. In Section~\ref{section:Sample} we present an overview of the HALOGAS sample; the observing configuration, reduction strategy and a discussion on our GBT beam model and how we convolve the WSRT data to avoid contamination from extended structure are outlined in Section~\ref{section:ObsRed}. The results from our comparison between the GBT and WSRT data for our initial four sources are discussed in Section~\ref{section:Results} with an investigation into how the diffuse $\hi$ environment relate to intrinsic galaxy properties following in Section~\ref{section:Discussion}. We then summarize our conclusions and commenting on future work in Section~\ref{section:conclusion}. | \label{section:conclusion} \begin{table*} \centering \resizebox{\textwidth}{!} {\begin{tabular}{lccccccc} \hline \hline Derived Properties & NGC891 & NGC925 & NGC2403 & NGC2997 & NGC4414 & NGC4565 & NGC6946 \Tstrut\Bstrut\\ \hline GBT $\hi$ Mass [10$^{9}$$\Msun$] & 3.86$\pm$0.19 & 5.79$\pm$0.29 & 3.39$\pm$0.37 & 7.0$\pm$1.0 & 5.43$\pm$0.27 & 7.33$\pm$0.37 & 3.80$\pm$0.69 \Tstrut\Bstrut\\ Conv. WSRT $\hi$ Mass [10$^{9}$$\Msun$] & 3.81$\pm$0.19 & 5.54$\pm$0.28 & --- & --- & 4.56$\pm$0.22 & 7.46$\pm$0.39 & --- \Tstrut\Bstrut\\ High-Res. WSRT $\hi$ Mass [10$^{9}$$\Msun$] & 3.90$\pm$0.18 & 5.57$\pm$0.28 & --- & --- & 4.63$\pm$0.23 & 7.55$\pm$0.38 & --- \Tstrut\Bstrut\\ M$_{*}$ [10$^{10}$$\Msun$] & 0.30$\pm$0.10 & 4.40$\pm$0.10 & 0.39$\pm$0.01 & 8.70$\pm$0.10 & 4.15$\pm$0.10 & 6.15$\pm$0.10 & 0.12$\pm$0.01 \Tstrut\Bstrut\\ M$_{bary}$ [10$^{10}$$\Msun$] & 0.82$\pm$0.02 & 5.20$\pm$0.10 & 0.86$\pm$0.03 & 9.71$\pm$0.18 & 4.89$\pm$0.10 & 7.15$\pm$0.10 & 6.29$\pm$0.59 \Tstrut\Bstrut\\ Deprojected Physical Area [kpc$^{2}$] & 603$\pm$6 & 490$\pm$6 & 122$\pm$3 & 726$\pm$11 & 241$\pm$24 & 976$\pm$10 & 416$\pm$3 \Tstrut\Bstrut\\ V$_{rot}$\tablenotemark{a} [km s$^{-1}$] & 212$\pm$2 & 105$\pm$2 & 196$\pm$1 & 145$\pm$3 & 218$\pm$6 & 244 & 315$\pm$12 \Tstrut\Bstrut\\ WSRT $f_{19}$ & 0.011$\pm$0.001 & 0.014$\pm$0.001 & --- & --- & 0.064$\pm$0.004 & 0.006$\pm$0.001 & 0.004$\pm$0.001 \Tstrut\Bstrut\\ GBT $f_{19}$ & 0.005$\pm$0.001 & 0.008$\pm$0.001 & 0.005$\pm$0.001 & 0.003$\pm$0.001 & 0.046$\pm$0.002 & 0.008$\pm$0.001 & 0.003$\pm$0.001 \Tstrut\Bstrut\\ \\[-1.0em] \hline \end{tabular}} \caption{Summary of Derived Properties} \tablenotetext{1}{Rotation velocity taken from \textit{HYPERLEDA} search} \label{tab:derivedParams} \end{table*} We presented an initial analysis of deep ($N_{HI}$$\sim$10$^{18}$ cm$^{-2}$) GBT observations of four sources (NGC891, NGC925, NGC4414, and NGC4565) out of 24 total sources in the HALOGAS survey. These observations are among the most sensitive $\hi$ observations of external galaxies to date. In order to directly compare interferometer and single dish data, we solve for an optimal smoothing kernel specific to each source and convolve the WST data to GBT angular resolution. Our main conclusions are: \begin{itemize} \item As we do not find significant spatially extended $\hi$ features, we conclude that the WSRT data do an excellent job recovering the diffuse (18 $\leq$ log$_{10}$($N_{HI}$) $\leq$ 19) $\hi$ around these four sources. In the case of NGC925, we detect about 20\% more $\hi$ than observations done with the VLA as part of the THINGS survey. The discrepancy is likely due in large part to the increased surface brightness sensitivity of the WSRT data since the ability to detect extended structure between the two surveys is very similar. The excellent agreement between the global $\hi$ profiles, cumulative $\hi$ mass as a function of $N_{HI}$, radial mean column density profiles, and radial cumulative flux for the GBT and convolved WSRT data provides additional evidence in support of this conclusion. \item The cumulative $\hi$ mass as a function of $\hi$ column density reveals the diffuse $\hi$ associated with these galaxies does not change significantly over the range $log_10\left(N_{HI}/cm^{-2}\right)$ = 18.0 to $log_{10}\left(N_{HI}/cm^{-2}\right)$ = 19.5. The flat behavior is consistent with predictions from simulations, which show the neutral fraction is around 1\% at $log_10\left(N_{HI}/cm^{-2}\right)$ = 18.0. Scaling our GBT beam model to the peak column density of the GBT data and repeating our analysis to essentially simulate an unresolved source produces a similarly flat distribution, which suggests the lowest column density bins include some values that trace the extended structure of the GBT beam. That said, there is generally a moderate offset between the data and model cumulative HI mass distributions. The overall agreement between the GBT and WSRT data sets, if extended to the other sources in our survey, suggests future surveys must probe column densities at the $\sim$10$^{17}$ cm$^{-2}$ level to increase the probability of detecting $\hi$ associated with cosmic web structure or possibly cold-mode accretion. \item We define a parameter, $f_{19}$, equal to one minus the ratio between the $\hi$ mass measured at and above log$_{10}$($N_{HI}$) = 19 and the source's total $\hi$ mass. We find that, on average (and excluding data that may suffer from resolution effects), this value is equal to 2\%, indicating the diffuse extended disks of these galaxies do not constitute a significant fraction of the overall mass. \end{itemize} One observational method to differentiate between inflow/outflow origins is a measure of metallicity using UV absorption lines (e.g., S II) utilizing the Cosmic Origins Spectrograph on the $\textit{Hubble}$ Space Telescope. If a significant diffuse $\hi$ feature is seen around a source as we analyze the full survey, and a fortuitous background quasar along the line of sight, a metallicity of $\sim$ 0.1 $Z_{\odot}$ would be highly indicative of a CGM origin. To establish or rule out cold mode accretion as a feasible avenue for nearby galaxies to refuel their gas content, we must continue to analyze galaxies within the HALOGAS sample that satisfy the mass constraints set by simulations, show large diffuse $\hi$ mass fractions and low SFRs, and reside in low density environments. Due to our small sample size we can only present foundational work to uncover any underlying correlations between large mass fractions of diffuse $\hi$ and galaxy properties. Future work will focus on the analysis techniques presented here in order investigate the origins of these large $\hi$ filaments, apply short spacing corrections to the WSRT data, and continue the investigation into role of $\hi$ in galaxy evolution. | 18 | 8 | 1808.02041 |
1808 | 1808.10712_arXiv.txt | \textcolor{black}{Knowledge of the atmospheric turbulence in the telescope line-of-sight is crucial for wide-field observations assisted by adaptive optics (AO)}, for which the \textcolor{black}{Point Spread Function} (PSF) becomes strongly elongated due to the anisoplanatism effect. This one must be modelled accurately to extrapolate the PSF anywhere across the Field of view (FOV) and improve the science exploitation. However, anisoplanatism is a function of the $\cnh$ profile, that is not directly accessible from single conjugate AO telemetry. \textcolor{black}{One may rely on external profilers, but recent studies have highlighted more than 10\% of discrepancies with AO internal measurements, while we aim better than 1\% of accuracy for PSF modelling.} To tackle this existing limitation, we present the Focal plane profiling (FPP), as a $\cnh$ profiling method that relies on post-AO focal plane images. We demonstrate such an approach complies with a 1\%-level of accuracy on the $\cnh$ estimation and establish how this accuracy varies regarding the calibration stars magnitudes and positions in the field. We highlight that photometry and astrometry errors due to PSF mis-modelling reaches respectively 1\% and 50$\mu$as using FPP on a Keck baseline, with a preliminary calibration using a star of magnitude H=14 at 20". We also validate this concept using Canada’s NRC-Herzberg HENOS testbed images in comparing FPP retrieval with alternative $\cnh$ measurements on HeNOS. The FPP approach allows to profile the $\cnh$ using the SCAO systems and improve significantly the PSF characterisation. Such a methodology is also ELT-size compliant and will be extrapolated to tomographic systems in a near future. | This paper focuses on improving the AO PSF characterisation in a wide-field by retrieving the distribution of atmospheric turbulence along altitude it depends on, referred as the $\cnh$ profile. We consider Single Conjugate Adaptive Optics (SCAO) system-assisted observations; \textcolor{black}{the correction provided by AO is optimal in the Guide star~(GS) direction, that can be either natural~(NGS) or artificial using laser~(LGS)}, but degrades across the field on account of the anisoplanatism effect~\cite{Fried1982}. This \textcolor{black}{latter} results from the spatial decorrelation of the incoming electric field phase that propagates through the atmosphere. The way this decorrelation occurs is a direct function of the $\cnh$ and outer scale profiles. As a consequence, anisoplanatism broaden the PSF and induces a spatial variation of the PSF morphology across the field.\\ However, the PSF model is one of the key limitations in the current exploitation of images of crowded-field stellar populations~\cite{Fritz2010,Yelda2010,Shodel2010} that are affected by anisoplanatism. To strengthen the data-processing outcome, we propose to improve the anisoplanatism characterisation. We have established a complete and general anisoplanatism model in~\cite{BeltramoMartin2018} as a function of the input $\cnh$ profile. This information is not accessible from the AO telemetry for SCAO systems, though. Dedicated instruments exist to monitor this profile~\cite{Osborn2015,Butterley2006,Wilson2002}, but they aim at characterising the observation site in terms of atmosphere quality and do not observe in the telescope line of sight. Consequently, their estimated profile deviate up to at least 10~\%~\cite{Ono2017}, from AO telemetry-based approaches available for multi-GS AO systems~\cite{Helin2018,Guesalaga2017,Laidlaw2016,Martin2016L3S,Neichel2014}. \textcolor{black}{However, according to \cite{BeltramoMartin2018}, 10\% of error on the $\cnh$ estimation may degrade the photometry and astrometry determination up to respectively 3\% and 300 $\mu$as in a 20"-FOV, while we are seeking to reach better than 1\% and 150 $\mu$as.} \\ We propose in this paper to rely on the SCAO-compensated PSFs available across the field. \textcolor{black}{We focus on SCAO systems that do not permit to identify the $\cnh$ from the telemetry and would benefit to have an internal images-based facility to retrieve the profile, either for post-processing or real-time application. The methodology we present can be extended to multi-GS systems, but at the cost of a larger numerical complexity to include the AO system control specificity, such as the tomographic reconstruction step or the optimal fitting in multi-conjugated AO. Before considering such systems, we will address the ground-layer AO case and compare our present approach to telemetry-based $\cnh$ estimation.}\\ Spatial variations of the AO PSF encodes the real $\cnh$ that affects images and the one that we want to determine. To extract the profile from images, we have developed the Focal plane profiling~(FPP) method as a non-linear least-square minimisation procedure that adjust a $\cnh$-dependent PSF model to match a collection of observations and deliver a joint estimation of the PSF model and $\cnh$ profile. If the model is not consistent with real atmosphere statistics, because we consider a profile over too few bins for instance, the minimisation process allows to mitigate these errors when extrapolating the PSF in the field. At the contrary, if we feed this inaccurate model with a wrong $\cnh$ without any feedback from a real PSF, we risk to amplify the error propagation and \textcolor{black}{degrade} the PSF extrapolation.\\ We describe the FPP algorithm in Sect.~\ref{S:FPP}. Sect.~\ref{S:simu} is dedicated to FPP performance assessment; we illustrate that the FPP allows to retrieve the $\cnh$ at a 1\%-level accuracy when bright stars are available. We present an application to PSF extrapolation on simulated images of NIRC2 \cite{McLean2000} at Keck II and evaluate what are the conditions in terms of calibration stars magnitudes and field location to decrease errors provoked by PSF-model indetermination on photometry and astrometry, down to respectively 1\% and 50$\mu$as. We make a step further in applying the FPP to Canada's NRC-Herzberg HENOS testbed images and compare successfully the FPP-retrieved $\cnh$ to existing measurements. | \label{S:conclusions} We present in this paper the focal plane profiling as a $\cnh$ retrieval method that relies on partially compensated AO images affected by anisoplanatism. It performs a non-linear least-square minimisation of a PSF model over observations and provide both a PSF model across the field and the $\cnh$ profile. To mitigate noise propagation and the sensitivity to unmodelled aberrations, such as field-dependent static aberrations, we must collect several PSF from the field; we show for the NIRC2 imaging camera at Keck II that the FPP can retrieve both atmosphere and PSF characteristics, photometry at 1\%-level accuracy and 50$\mu$as astrometry if we get $n_\text{psf}$ PSFs of magnitude given by $H=14 + 2.5\times\log(n_\text{psf}\times T_\text{exp}/30)$, with a breakdown to H=15.5 mag when picking off 4 stars in the field. We have deployed this approach on the HeNOS testbench where $\cnh$ values are measured from WFS cross-correlation. Thanks to FPP, we retrieved a profile that complies with WFS-based measurements when using three stars distributed over 4.5" with $\theta_0 = 0.854"$. We demonstrate that collecting more stars allows to mitigate model errors such as field-dependent static aberrations for instance.\\ We focused in this paper for the $\cnh$ profiling in the purpose of assessing the reliability and limitations of this method. Our next work will consist in deploying both FPP and PSF-reconstruction technique to asses what are the potential gains on crowded fields observations, such as the Galactic center with Keck and the ELT with MICADO. Furthermore, we will investigate to extend the FPP to tomographic systems for improving the $\cnh$ profiling, especially in the purpose to deploy such an approach to the multi-conjugated system MAORY coupled with MICADO, or the laser-tomographic mode of HARMONI and METIS.\\ | 18 | 8 | 1808.10712 |
1808 | 1808.04101_arXiv.txt | A Galactic supernova (SN) axion signal would be detected in a future neutrino Mton-class water Cherenkov detector, such as the proposed Hyper-Kamiokande in Japan. The main detection channel for axions is absorption on the oxygen nuclei in the water. The subsequent oxygen de-excitation leads to a potentially detectable gamma signal. In this contribution we present a calculation of the SN axion signal and discuss its detectability in Hyper-Kamiokande. | Low-mass axions can be copiously produced in a core-collapse supernova (SN) affecting the neutrino burst. In fact, bounds on axions have been placed studying the neutrino signal from SN 1987A. In particular, in \cite{engels} it was pointed out that a SN axion burst could produce an observable signal in a neutrino water Cherenkov detector by oxygen absorption. The following oxygen de-excitation would produce a photon signal. Currently it has been proposed a future Mton-class neutrino water Cherenkov detector, Hyper-Kamiokande, in Japan. Motivated by this exciting situation, we find it worthwhile to take a fresh look at the possibility of detecting a SN axion burst. In this contribution we present preliminary results of an updated calculation of the SN axion signal in Hyper-Kamiokande. We refer the interested reader to \cite{master,prep} for further details. | We evaluated the Hyper-Kamiokande potential to detect the axion burst associated with a Galactic SN event. We found that axions in the free-streaming regime would be potentially detectable if a careful reduction of the neutrino background is performed. On the other hand, in the trapping regime the axion signal would dominate over the neutrino one, being easily detectable. Therefore a Galactic SN explosion would be a once in a lifetime opportunity for detecting axions. | 18 | 8 | 1808.04101 |
1808 | 1808.06618_arXiv.txt | Fast flavor conversions of supernova neutrinos, possible near the neutrinosphere, depends on an interesting interplay of collisions and neutrino oscillations. Contrary to naive expectations, the rate of self-induced neutrino oscillations, due to neutrino-neutrino forward scattering, comfortably exceeds the rate of collisions even deep inside the supernova core. Consistently accounting for collisions and oscillations, we present the first calculations to show that collisions can create the conditions for fast flavor conversions of neutrinos, but oscillations can continue without significant damping thereafter. This may have interesting consequences for supernova explosions and the nature of its associated neutrino emission. | $} \label{sec:appen} Assuming a neutrino energy $E_\nu$ and a distance from the supernova centre $r$, the inverse of the scattering rate for the process $\nu_e+n\rightarrow e^-+p$ ($\bar{\nu}_e+p\rightarrow e^+n$) is given by \cite{Bruenn:1985en} \begin{equation} \Gamma^{-1}_{\barparenb{\nu}_e}(E_\nu,r)={\frac{G^2}{\pi}n_{n\,(p)}(g_V^2+3g_A^2)(E_\nu+Q)^2\left(1-\frac{m_e^2}{(E_{\nu}+Q)^2}\right)}\,, \label{collision_rate_nue} \end{equation} where $m_e$ is the mass of the electron, $G^2=5.18\times10^{-44}$ MeV$^{-2}$ cm$^2$, $Q=1.2935$ MeV, $g_V=1$, $g_A=1.23$ and $n_{n\,(p)}(r)$ is the neutron (proton) density at a distance $r$. In Eq.\,\ref{collision_rate_nue} we are neglecting the decrease in the scattering rate due to nucleon and electron final state blocking. Such an effect is relevant when temperatures are lower than degeneracy, e.g., when matter density becomes of the order of the nuclear density. However, since our purpose is a comparison with $\mu$ (see Fig.~1), we show that $\mu$ dominates over $\Gamma$ even with such assumptions. For each $r$ we calculate an average of the scattering rate using the neutrino density as weight function \begin{equation} \langle {\Gamma}_{\barparenb{\nu}_e}(r) \rangle =\frac{\int_0^\infty dE_\nu\Gamma_{\barparenb{\nu}_e}(E_\nu,r)n_{\barparenb{\nu}_e}}{\int_0^\infty dE_\nu n_{\barparenb{\nu}_e}}\,. \label{collision_rate_nue_average} \end{equation} A similar average is performed also in the calculation of the neutrino potential $\mu(r)$. These energy-averaged quantities are shown in the main text in Fig.~\ref{fig:profiles}, where we omit the notation $\langle \cdots \rangle$ for typographic clarity. | 18 | 8 | 1808.06618 |
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1808 | 1808.05794_arXiv.txt | { In the dawn of the multi-messenger era including gravitational waves, which was marked by the first ever coincident detection of gravitational waves and electromagnetic radiation it is important to lay back and think about established knowledge. Numerical simulations of binary neutron star mergers and simulations of short GRB jets have to combine efforts in order to understand such complicated and phenomenologicaly rich explosions. We review the status of numerical relativity simulations with respect to any jet or magnetized outflow produced after merger. We compare what is known from such simulations, with what is used and obtained from short GRB jet simulations propagating through the BNS ejecta. We point out facts that are established and can be considered known, and things that need to be further revised and/or clarified. } \keyword{keyword 1; keyword 2; keyword 3 (list three to ten pertinent keywords specific to the article, yet reasonably common within the subject discipline.)} \begin{document} \setcounter{section}{-1} % | The detection of GW170817 marked the dawn of the multi-messenger gravitational-wave era \citep{Abbott2017, Abbott2017b}. The subsequent observation of a short gamma-ray burst (GRB) almost $\sim 1.7$ seconds after merger \citep{Goldstein2017,Savchenko2017}, showed that a least a subset of short GRBs is produced by binary neutron star (BNS) mergers. Hours after merger a precise localization could be established through optical observations of GW170817 \citep{Coulter2017,Soares-Santos2017}, identifying the host as galaxy NGC 4993, which is at a distance of 40 megaparsecs. Further detection in UV/optical/Infrared established perennially the connection of BNS mergers with a kilonova (macronova) \citep{Soares-Santos2017,Arcavi2017,Nicholl2017,Pian2017, Smartt2017,Tanvir2017,Ustumi2017, Kilpatrick2017,Kasliwal2017,Covino2017,Cowperthwaite2017, Buckley2018,Drout2017,Evans2017,Arcavi2018}. A coincident detection of GW and a short GRB from a BNS merger was long ago conjectured that short-duration GRBs come from BNS mergers \citep{Eichler89,Narayan92,Mochkovitch93}. These unprecedented observations open new windows and insights for the in depth study of such objects and such events. The observation opened the possibility of constraining the maximum mass of neutron stars and the equation of state (EOS) \citep{Annala2017,Bauswein2017b,Margalit2017,Radice2017b, Rezzolla2017,Ruiz2017,Shibata2017c,De2018,Tews2018a,Tews2018,Alsing2017, Burgio2018,Raithel2018,Paschalidis2017}. It was long proposed that a BNS merger would give rise to emission powered by the radioactive decay of r-process nuclei \citep{Lattimer74,Li:1998}. Several groups concluded that this was the case for the optical/NIR emission that followed GW170817 \citep{Kasen2017,Drout2017,Tanaka2017,Kasliwal2017,Murguia-Berthier2017, Waxman2017,Villar2017, Tanvir2017, Utsumi2017,Perego2017,Metzger2018}. This observation triggered further modeling for the actual components that give rise to this emission and how these components were produced. The prompt gamma-ray emission was reported in \citep{Goldstein2017,Savchenko2017}. It was the most faint (short or long) GRB ever detected \citep{Goldstein2017}. The first X-ray afterglow observations came nine days after merger \citep{Evans2017,Troja2017,Margutti2017}. The first radio counterparts came later, sixteen days after merger \citep{Hallinan2017,Alexander2017}. Every information that would come from the afterglow observations would be invaluable to reveal the nature of the outflow and its structure. A relativistic outflow from a BNS merger was indeed observed \citep{Alexander2017,Haggard2017}. Was that the most peculiar short GRB ever detected \citep{Kasliwal2017,Granot2017}? The continuous rising of the afterglow the first 100 days suggested that a simple top-hat\footnote{A top-hat jet is one with constant Lorentz factor and emissivity within the jet that goes sharply to zero outside of jet opening angle. It is the simplest model to explain GRBs have been widely used to explain GRB properties.} jet model seen off-axis is not adequate for explanation\cite{Alexander2017,Mooley2018,Ruan2018,Pooley2018, Margutti2018,Lyman2018}. However, at a 100 days after merger the data could not exclude other jet structure or cocoon models. Energy injection was evident at that time \citep{Li2018}. Then, a turnover in the light curve appeared close to 200 days \citep{Dobie2018,Alexander2018,Nynka2018}. This emission is well understood and comes from the interaction of the outflow as it smashes into the inter-stellar medium producing a shock which accelerates electrons that radiate synchrotron radiation and can give a great insight in the whole structure of the initial outflow. In order to digest all these new insightful observations, and the yet to come in the next years we have to combine all pieces available. What has been achieved from BNS numerical relativity simulations has to be part of any adequate modeling of short GRB outflows. These are: the ejected matter and the production of neutrino driven winds, the enormous magnetic field evolved in the merger process, and its amplification during merger, the actual possibility of launching a relativistic outflow after merger are the starting points given by numerical relativity. Is it a stable magnetar or the collapse to a black hole (BH) torus system that powers an outflow? In what follows we try to present results from numerical relativity BNS simulations relevant for short GRBs. Afterwards, we turn our attention to efforts in short GRB jet simulations propagating through the BNS ejecta. This is a rather focused review on what we know from numerical relativity concerning short GRBs and and how this knowledge is applied to short GRB simulations. It will not at all follow the path of excellent reviews that exist in the subject of BNS mergers. For the interested reader we cite several detailed reviews of subjects relevant to the detection of a BNS merger, a short GRB and a kilonova. Detailed reviews of all the aspects of numerical relativity and its applications to BNS mergers \citep{Faber2012:lrr,Baiotti2016}, a focused review on BH - neutron star binaries \citep{Shibata06c}, a review discussing the connection between BNS mergers and short GRBs in numerical relativity results \citep{Paschalidis2016}, observational aspects of short GRBs and connection to BNS mergers \citep{Berger2013b,Fong2015}, regarding BNS merger and electromagnetic counterparts from kilonova \citep{Rosswog2015,Fernandez2015b,Thielemann2017b,Metzger2017}, a review on rotating stars in relativity with applications on the post merger phase \citep{Paschalidis2017b}, reviews for short GRBs \citep{Nakar:2007yr,Lee:2007js} and detailed reviews on all aspects of GRBs \citep{Piran:2004ba,Meszaros:2006rc,Kumar2014}. In section \ref{sec:BNS} we review the relevant knowledge from BNS simulations. We mainly follow results from magnetohydrodynamic (MHD) simulations in BNS studies. At the end of section \ref{sec:BNS} we show the different paths that a BNS may follow after merger with respect to the achieved magnetic energy growth during merger. This translates to the total mass of the binary. In section \ref{sec:shortGRB} we follow the studies that focus on the interaction of a BNS relativistic outflow passing through the matter that has been ejected during merger. In section \ref{sec:conc} we conclude. | \label{sec:conc} In the years to come more detections of BNS mergers are expected from ground-based interferometers. It is important to analyze in detail observations of GW170817 and all its EM counterparts, starting with GRB170817A. It is equally important to reproduce realistic physics through numerical simulations in order to match and explain observations. This brief review can act as a quick introduction to BNS numerical relativity simulations for people interested in short GRB outflows through BNS ejecta. Or as a brief introduction to short GRB jet simulations and setups used for people from BNS merger simulations. Overall, we want to point out the importance of combining knowledge from both paths in order for a consistent picture to be drawn at the end. In section \ref{sec:BNS} we went through studies from numerical relativity for magnetized BNS mergers. We highlight important aspects of this physical process as given in the literature. Points, such as the magnetic field amplification, the difficulty of launching a relativistic jet and certainly the mass ejection during merger and all possible winds produced after merger can become clear through detailed studies. At the end of the section we state several important points (importance is a subjective criterion). Next step is to take these ingredients from BNS simulations and study any outflow emerging after merger. A relativistic outflow has been observed from a BNS merger \citep{Alexander2017,Haggard2017}. Thus, we need to understand how it was launched and what is the initial structure of this outflow. Furthermore, how it will evolve through its interaction with the BNS ejecta. In section \ref{sec:shortGRB} we briefly go through previous works on these aspects. This is a rapidly evolving sub-field, especially after the detection. Now, any model and idea can be simulated and be exposed to the data that followed GW170817. However, we should keep in mind that a BNS can have a different evolution even with a slightly difference in mass. At the end, modeling and studying outflows of such events should be inspired by GW170817. \vspace{6pt} | 18 | 8 | 1808.05794 |
1808 | 1808.07875_arXiv.txt | We present the discovery of ASASSN-18ey (MAXI J1820+070), a new black hole low-mass X-ray binary discovered by the All-Sky Automated Survey for SuperNovae (ASAS-SN). A week after ASAS-SN discovered ASASSN-18ey as an optical transient, it was detected as an X-ray transient by MAXI/GCS. Here, we analyze ASAS-SN and Asteroid Terrestrial-impact Last Alert System (ATLAS) pre-outburst optical light curves, finding evidence of intrinsic variability for several years prior to the outburst. While there was no long-term rise leading to outburst, as has been seen in several other systems, the start of the outburst in the optical preceded that in the X-rays by $7.20\pm0.97~\rm days$. We analyze the spectroscopic evolution of ASASSN-18ey from pre-maximum to $> 100~\rm days$ post-maximum. The spectra of ASASSN-18ey exhibit broad, asymmetric, double-peaked H$\alpha$ emission. The Bowen blend ($\lambda \approx 4650$\AA) in the post-maximum spectra shows highly variable double-peaked profiles, likely arising from irradiation of the companion by the accretion disk, typical of low-mass X-ray binaries. The optical and X-ray luminosities of ASASSN-18ey are consistent with black hole low-mass X-ray binaries, both in outburst and quiescence. | \label{sec:intro} Low-mass X-ray binaries (LMXBs) consist of compact objects, either a neutron star (NS) or a black hole (BH), accreting material from a donor star with a typical mass of $M_{\rm{donor}} \lesssim 1~M_\odot$. The compact object is surrounded by an accretion disk fed by a donor star undergoing Roche Lobe overflow (RLOF). Observationally, LMXBs can be classified as either transient/outbursting sources or persistent/non-outbursting sources, with the caveat that transient LMXBs can go undetected for years or decades while in quiescence since their X-ray luminosity is low ($L_X\sim 10^{32}~\rm{erg}\;\rm s^{-1})$. During an outburst, these systems increase by several orders of magnitude in both X-ray and optical luminosity, routinely leading to their discovery. Conversely, persistent X-ray binaries have higher continuous X-ray luminosities ($L_X\sim 10^{36-38}~\rm{erg}\;\rm s^{-1}$), and the majority of these sources have been discovered by all-sky X-ray surveys. Transient LMXBs can be grouped into one of three categories based on their X-ray spectral state: high/soft/thermal, low/hard and very high/steep power law \citep[see ][for an overview of BH LMXB X-ray properties]{remillard06}. The high/soft state is dominated by thermal disk emission, with little to no power law component. Conversely, the low/hard state is dominated by the power law emission, contributing $\gtrsim 80\%$ of the observed flux. Finally, the very high state is characterized by a steep power law ($\Gamma > 2$) and usually dominates the X-ray spectra when BH LMXBs approach the Eddington limit. Throughout a single outburst, a LMXB usually experiences at least two of these states as the accretion rate onto the compact object evolves with time. These transient LMXB outbursts are likely driven by thermal and viscous instabilities in the accretion disk \citep{dubus01, lasota01}. Both NS and BH LMXBs can be transient sources, although their outbursts are quantitatively different (see \citealp{done07} for a review). NS LMXBs generally host smaller accretion disks than their BH counterparts due to tidal truncation by the donor and the lower mass of the compact object. The smaller accretion disk is less likely to be unstable for a given donor star, and even when unstable, these smaller disks result in lower amplitude outbursts than those seen in BH LMXBs \citep{done07}. BH LMXBs can exhibit optical outbursts of $\gtrsim 5$~mags \citep{corral15}, and go years between consecutive outbursts \citep[e.g., ][]{russell18}. These outbursting LMXBs are well-described by the classical disk instability model (DIM) with modifications to account for self-irradiation \citep[DIM+irradiation, ][]{dubus01}. \name was discovered by the All-Sky Automated Survey for SuperNovae (ASAS-SN, see \citealp{shappee14} and \citealp{kochanek17} for details on cameras, filters, and zero-points) on UT 2018-03-06.58 (MJD 58184.079861) at $\textrm{RA}=18^{\rm h}20^{\rm m}21.\!\!^{\rm{s}}9$ $\textrm{Dec.}=+07\degree11'07.\!\!''3$ (J2000) with a $V$-band magnitude of 14.88. It was publicly released within hours of discovery.\footnote{\url{http://www.astronomy.ohio-state.edu/asassn/transients.html}} The source was undetected ($V > 16.7$ mag) on UT 2018-03-02.59, roughly 4 days prior. Because of its coincidence with a $G=17.8$ mag Gaia source \citep[ID 4477902563164690816, ][]{gaia1}, ASASSN-18ey was initially labeled as a cataclysmic variable (CV) candidate. Then, six days later on 2018-03-11, the Monitor of All-sky X-ray Image \citep[MAXI, ][]{MAXIref} Gas Slit Camera \citep[GSC, ][]{GSCref} nova alert system detected a bright X-ray transient at the same location (MAXI J1820+070, $32 \pm 9~\rm{mCrab}$, $4-10~\rm{keV}$, \citealp{ATel11399,ATel11400}). \begin{figure*} \centering \includegraphics[width=0.9\linewidth]{PreOutburst.pdf} \caption{Pre-outburst light curve of \name from ASAS-SN and ATLAS. Green diamonds represent ASAS-SN $V$-band, blue squares ATLAS $c$-band, and orange squares ATLAS $o$-band observations. There are obvious correlations between observations conducted in different filters, indicating much of the variability is intrinsic to \name and not nightly scatter. Fluxes are \textit{not} corrected for interstellar extinction.} \label{fig:pre-outburst} \end{figure*} Several teams carried out follow-up observations after the MAXI transient alert due to the intrinsic brightness and unknown nature of \name. Follow-up X-ray observations of \name were conducted with NICER \citep{ATel11423, ATel11576, ATel11823}, INTEGRAL \citep{ATel11478, ATel11488, ATel11490}, and XRT/BAT \citep{ATel11403, ATel11427, ATel11578}. The first suggestion of \name being a BH LMXB and the detection of a possible state transition were posted by \citet{ATel11418} and \citet{ATel11820}, respectively. Optical spectra showed signatures typically associated with LMXBs in outburst including \ion{He}{1} and \ion{He}{2} in emission combined with an evolving H$\alpha$ emission profile \citep{ATel11424, ATel11425, ATel11480, ATel11481}. Subsequent optical observations revealed an optical period of $\sim 3.4~\rm{hr}$ \citep{ATel11596}, correlations between X-ray and optical brightness \citep{ATel11432, ATel11510, ATel11574}, linear polarization \citep{ATel11445}, and \newline(sub-)second flaring \citep{ATel11421, ATel11426, ATel11437,ATel11510, ATel11591, ATel11723, ATel11824}. These rapid photometric variations were confirmed in the near-infrared by \citet{ATel11451}, and a large IR excess compared to archival 2MASS data was noted by \citet{ATel11458} and \citet{ATel11855}. Radio observations may have detected a forming jet and its ensuing quenching \citep{ATel11420, ATel11439, ATel11440, ATel11539, ATel11540, ATel11609, ATel11827, ATel11831, ATel11887}. Using the \textit{Gaia} DR2 parallax \citep{gaia1,gaia2,gaia3} and \citet{bailerjones18} distance priors, \name is located at a distance of $d = 3.06 ^{+1.54} _{-0.82}~\rm{kpc}$, with a total reddening of $E(B-V)=0.197$ mag \citep{SF11}, corresponding to $A_V = 0.614$ mag assuming $R_V=3.1$. We discuss pre-outburst data in \S\ref{sec:preoutburst}, analyze the rising light curves in \S\ref{sec:lc}, and qualitatively examine the pre- and post-maximum spectra in \S\ref{sec:spec_evol}. Finally, in \S\ref{sec:new_BHXB}, we discuss our conclusion that \name is likely a new BH LMXB in outburst. | 18 | 8 | 1808.07875 |
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1808 | 1808.09435_arXiv.txt | Many problems in Galactic structure benefit from accurate knowledge of the distance to the Galactic center, $R_0$. This crucial value has recently been measured to unprecedented accuracy --- $R_0 = 8.122 \pm 0.031\;\mathrm{kpc}$ --- thanks to relativistic effects observed during the pericenter passage of a star orbiting the central supermassive black hole \citep{GRAVITY}. Combined with the observed proper motion of Sgr A$^*$ \citep{SgrAstar}, this Galactocentric distance implies a transverse speed of the sun of $245.58 \pm 1.32\;\kms$. Adopting $V_{\sun} = 12.24 \pm 0.47\;\kms$ \citep{solarmotion} for the solar motion leads to a circular speed of $\Theta_0 = 233.34 \pm 1.40\;\kms$ for the Local Standard of Rest. In light of this development, I provide here an update to the Milky Way model of \citet{M16}. In addition to the new Galactocentric distance, new measurements of the terminal velocities in the first quadrant are also available \citep{MGDQ1}. Combining these with fourth quadrant terminal velocities \citep{MGDQ4} and the Galactic constants ($R_0$, $\Theta_0$) = (8.122 kpc, 233.3 \kms) provides a remarkably detailed picture of the Galactic rotation curve. The method of \citet{M16} applies the Radial Acceleration Relation \citep[RAR:][]{RAR} to derive the azimuthally averaged baryonic surface density profile $\Sigma(R)$ from the Galactic rotation curve. Features in $V(R)$ have corresponding features in $\Sigma(R)$ that map well to independently known features like the Centaurus spiral arm. The resulting $\Sigma(R)$ departs from a pure exponential profile in a way that is typical of other spiral galaxies. The model Q4MB of \citet{M16} provides the starting point here as its bulge model provides an excellent match to the inner rotation curve obtained from the detailed modeling of \citet{Portail2015} when both are scaled to the same $R_0$. The pattern of bumps and wiggles at larger radii --- presumably the signature of massive spiral arms like Centaurus --- is a good match to the terminal velocity data in both quadrants. The new, slightly larger $R_0$ makes the Milky Way more massive, a net increase of 9\% over the Q4MB model, with correspondingly higher surface densities. | 18 | 8 | 1808.09435 |
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1808 | 1808.07320_arXiv.txt | This is a guide for preparing papers for \textit{Monthly Notices of the Royal Astronomical Society} using the \verb'mnras' \LaTeX\ package. It provides instructions for using the additional features in the document class. This is not a general guide on how to use \LaTeX, and nor does it replace the journal's instructions to authors. See \texttt{mnras\_template.tex} for a simple template. The abstract of the paper. | The journal \textit{Monthly Notices of the Royal Astronomical Society} (MNRAS) encourages authors to prepare their papers using \LaTeX. The style file \verb'mnras.cls' can be used to approximate the final appearance of the journal, and provides numerous features to simplify the preparation of papers. This document, \verb'mnras_guide.tex', provides guidance on using that style file and the features it enables. This is not a general guide on how to use \LaTeX, of which many excellent examples already exist. We particularly recommend \textit{Wikibooks \LaTeX}\footnote{\url{https://en.wikibooks.org/wiki/LaTeX}}, a collaborative online textbook which is of use to both beginners and experts. Alternatively there are several other online resources, and most academic libraries also hold suitable beginner's guides. For guidance on the contents of papers, journal style, and how to submit a paper, see the MNRAS Instructions to Authors\footnote{\label{foot:itas}\url{http://www.oxfordjournals.org/our_journals/mnras/for_authors/}}. Only technical issues with the \LaTeX\ class are considered here. Sections and lists are generally the same as in the standard \LaTeX\ classes. \subsection{Sections} \label{sec:sections} Sections are entered in the usual way, using \verb' | 18 | 8 | 1808.07320 |
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1808 | 1808.00670_arXiv.txt | Cosmic strings and primordial black holes (PBHs) commonly and naturally form in many scenarios describing the early universe. Here we show that if both cosmic strings and PBHs are present, their interaction leads to a range of interesting consequences. At the time of their formation, the PBHs get attached to the strings and influence their evolution, leading to the formation of black-hole-string networks and commonly to the suppression of loop production in a range of redshifts. Subsequently, reconnections within the network give rise to small nets made of several black holes and connecting strings. The number of black holes in the network as well as the stability of the nets depend on the topological properties of the strings. The nets oscillate and shrink exponentially due to the emission of gravitational waves. This leads to potentially observable string-driven mergers of PBHs. The strings can keep PBHs from galactic halos, making the current bounds on PBHs not generally applicable. Alternatively, heavy PBHs can drag low-tension strings into the centers of galaxies. The superconducting strings can appear as radio filaments pointing towards supermassive black holes. | Primordial black holes (PBHs) could be formed early in the radiation era and could have played an important role in the evolution of the universe (see Refs.~\cite{Carr:2009jm,Sasaki:2018dmp} for recent reviews). They could serve as seeds for supermassive black holes observed at the centers of most galaxies and could account for merging BH binaries currently observed by LIGO. It has also been suggested that PBHs could constitute a substantial part, if not all, of dark matter. A seemingly unrelated direction of cosmological research is the formation and evolution of topological defects, in particular cosmic strings. Strings are predicted in a wide class of particle physics models and can produce a variety of astrophysical effects (see \cite{book,Copeland} for a review). Some models predict more complicated defects -- monopoles connecting by strings, with $N\geq 2$ strings attached to each monopole. In the case of $N=2$ the defects are called "necklaces", with monopoles and antimonopoles alternating like beads along the string. A discovery of cosmic strings would open a window into physics of very high energies, and new observational bounds on strings may rule out classes of particle physics models. In this paper we discuss what happens when both strings and PBHs are present in the universe. We find that they generically form a network where PBHs are connected by strings. The evolution and observational effects of strings may then be significantly affected by PBHs and vice versa. In particular, observational bounds on the string energy scale are strongly enhanced in some scenarios. Another interesting effect is the increased merger rate for PBHs connected by strings. | In this work we introduced a new class of models in which cosmic strings and PBHs are both present. The interaction between the two is dramatic and results in the formation of an infinite string network with BHs at the nodes. In the course of the following evolution, strings cross and reconnect, and some fraction of BHs is detached from the network. We have discussed two limiting scenarios, providing upper and lower bounds on the density of BHs that remain in the infinite network. A more definitive analysis of network evolution will require numerical simulations. One of the main observational signatures of cosmic strings is the gravitational wave (GW) background produced by oscillating string loops. In the presence of PBHs, the loop production by the network is significantly modified, resulting in a modified GW spectrum. In addition, the network provides another source of GW -- small oscillating nets of black holes and strings. The nets lose their energy by gravitational radiation and shrink exponentially on a timescale of $ \tau_{GW}\sim t_H (M/100M_{\odot})(10^{-10}/G\mu)^{-2}$, where $t_H$ is the current age of the Universe. If such objects exist and are detected by LIGO, Pulsar Timing Arrays, or LISA, this would give most valuable insights for particle physics and astrophysics alike. To further investigate observational signatures of the model, the properties of nets with black holes need to be explored in greater detail. | 18 | 8 | 1808.00670 |
1808 | 1808.08757_arXiv.txt | The detection of bright X-ray features and large TeV halos around old pulsars that have escaped their parent Supernova Remnants and are interacting directly with the ISM, suggest that high energy particles, more likely high energy pairs, can escape from these systems, and that this escape if far more complex than a simple diffusive model can predict. Here we present for the first time a detailed analysis of how high energy particles escape from the head of the bow shock. In particular we focus our attention on the role of the magnetic field geometry, and the inclination of the pulsar spin axis with respect to the direction of the pulsar kick velocity. We show that asymmetries in the escape pattern of charged particles are common, and they are strongly energy dependent. More interestingly we show that the flow of particles from bow-shock pulsar wind nebulae is likely to be charge separated, which might have profound consequences on the way such flow interacts with the ISM magnetic field, driving local turbulence. | \label{sec:intro} Once escaped from their parent Supernova Remnant (SNR), usually a few tens of thousands of years after the Supernova explosion, pulsars (PSRs) begin to interact directly with the ISM \citep{Gaensler_Slane06a,Bucciantini_Bandiera01a}. The ram pressure balance between the pulsar wind and the supersonic ISM flow (in the reference frame of the PSR itself) leads to the formation of a cometary nebula known as bow-shock pulsar wind nebula (BSPWN) \citep{Wilkin96a,Bucciantini_Bandiera01a,Bucciantini02b}. These nebulae have been observed in H$_{\alpha}$ emission, due to neutral Hydrogen of the ISM interacting through charge-exchange with the ionized component of the ISM itself, shocked in the outer bow-shock \citep{Kulkarni_Hester88a,Cordes_Romani+93a,Bell_Bailes+95a,van-Kerkwijk_Kulkarni01a,Jones_Stappers+02a,Brownsberger_Romani14a,Romani_Slane+17a}, and in non-thermal radio/X-rays associated with the shocked pulsar wind, flowing backward into long tails \citep{Arzoumanian_Cordes+04a,Kargaltsev_Pavlov+17a,Kargaltsev_Misanovic+08a,Gaensler_van-der-Swaluw+04a,Yusef-Zadeh_Gaensler05a,Li_Lu+05a,Gaensler05a,Chatterjee_Gaensler+05a,Ng_Camilo+09a,Hales_Gaensler+09a,Ng_Gaensler+10a,De-Luca_Marelli+11a,Marelli_De-Luca+13a,Jakobsen_Tomsick+14a,Misanovic_Pavlov+08a,Posselt_Pavlov+17a,Klingler_Rangelov+16a,Ng_Bucciantini+12a}. Modeling of BSPWNe have progressed in the last decades from simple classical axisymmetric hydrodynamics \citep{Bucciantini02b} to relativistic MHD \citep{Bucciantini_Amato+05a} and full and/or simplified 3D \citep{Vigelius_Melatos+07a,Bucciantini18b,Barkov_Lyutikov18a}. All of those models are based on the assumption that the pulsar wind behaves as a fluid, and as such is fully confined within a contact discontinuity (CD) bounding the non-thermal tail. However, old pulsars are often found embedded into extended TeV halos \citep{Gallant07a,Helfand_Gotthelf+07a,H.E.S.S.Collaboration_Abramowski+14a,H.E.S.S.Collaboration_Abdalla+18a,Abeysekara_Albert+17a}, that do not resemble at all the cometary structures found in simulations. It has been suggested that those PSRs might be slowly moving and still confined into their parent SNR. However more recently an extended TeV halo was found surrounding Geminga \citep{Abeysekara_Albert+17a}, which is a well known BSPWN \citep{Posselt_Pavlov+17a}, where X-rays clearly show a cometary shape. On the other hand, in the Guitar nebula associated with PSR B2224+65 \citep{Hui_Becker07a} and in the Lighthouse nebula associated to PSR J1101-6103 \citep{Pavan_Bordas+14a}, bright X-ray features clearly protruding out of the system, almost orthogonal to the PSR motion, are observed, that contrast strongly with what fluid models predict. Previous studies have attempted to explain the presence of one-sided X-ray features invoking a preferential collimated escape of high energy particles from the PWN \citep{Bandiera08a}. Such directed outflows have been associated with localized reconnection at the magnetopause (the CD) between the magnetic field line of the PWN, and those of the ISM, in a fashion not dissimilar to what is known to happen on the dayside Earth magnetopause, where it interacts with the solar wind \citep{Scholer03a,Frey_Phan+03a,Faganello_Califano+12a,Fuselier_Trattner+12a,Fuselier_Frahm+14b}. However reconnection at the magnetopause is known to be patchy and sporadic and to lead only to marginal flux transfer \citep{Kan88a,Pinnock_Rodger+95a,Fear_Trenchi+17a} in contrast with the persistence of those X-ray features, and on the large energetics (comparable to the PSR spindown) required to power them. Moreover, reconnection is likely to affect only those low energy particles that are bound to follow magnetic field lines, while existing models invoke the presence of particles with typical Larmor radii comparable with the size of the bow-shock, that are likely unable to feel the existence of small reconnecting regions. The problem of the escape of particles from old BSPWNe, is also relevant in the context of dark matter searches. Pulsars are likely the most efficient antimatter factories in the Galaxy, and BSPWNe have could be one of the major, if not the main, contributor to the positron excess observed by PAMELA \citep{Adriani_Barbarino+09a,Hooper_Blasi+09a,Blasi_Amato11a,Adriani_Bazilevskaya+13a,Aguilar_Alberti+13a}, in competition with dark matter \citep{Wang_Pun+06a}. Of course the flow dynamics in these systems can be quite complex, and many key factors (turbulence, differential acceleration at the termination shock, etc...) can lead to anisotropies in the escape of high energy particles. Here, however, our interest is focused in singling out the specific role of the magnetic field geometry, using a simplified model for the magnetic field structure in the head of these nebulae. We trace the trajectories of charged particles that emerge out of the PWN in its head, and assess the level of anisotropy associated to the emerging flow. In Sect.~\ref{sec:traj} we describe how we model the magnetic field structure and particle trajectories. In Sect.~\ref{sec:res} we illustrate our results for different magnetic field configurations and particle rigidities. In Sect.~\ref{sec:conc} we state our conclusions. | \label{sec:conc} Using a simplified model for the flow geometry and magnetic field structure in the head of a bow-shock pulsar wind nebula, we have investigated the escape of high energy charged particles, and in particular the role of the magnetic field in creating anisotropies in the escaping flow. Following their trajectories in the electric and magnetic field, we have try to assess how much the escape probability depends on the magnetic field geometry, in terms of strength of the field (the particle effective rigidity), the relative inclination of the pulsar spin axis and pulsar kick velocity, and in particular we focus on various kinds of anisotropies: {\it head to tail}, {\it inclination}, and especially {\it charge} anisotropy. We want to stress here that there are several other possible sources of asymmetry in the escape of particles from BSPWNe: the energy flow in the wind can have a strong latitudinal dependence \citep{Spitkovsky06a,Tchekhovskoy_Philippov+16a} which will affect the shape of the termination shock, and the structure of the flow downstream of it, leading to the formation of fast channels \citep{Del-Zanna_Amato+04a}; particles acceleration at the shock can depend strongly on the local conditions, in term of inclination and magnetization \citep{Spitkovsky08a,Sironi_Spitkovsky09a,Sironi_Spitkovsky11a}, such that the shock itself could introduce major asymmetries in the way particle are in injected; shear at the CD between the internal pulsar wind flow moving at a speed $\sim 0.3-0.6 c$, and the slow moving outer shocked ISM (with velocity comparable to the pulsar kick velocity), can produce Kelvin-Helmholtz instability which might disrupt the magnetopause that is found at the CD in laminar models, and introduce a further source of scattering and turbulence inside the nebula. This will likely tend to isotropize the escape. All of these effects can alter substantially the resulting escape pattern, and they can also introduce major time dependencies. However, even if the pattern of escape particles can differ substantially from what we found, there are a few results from our study that we deem robust: \begin{itemize} \item The fraction of escaping particles and the escape geometry is strongly energy dependent, and changing the energy from $\mathcal{R} =0.1$ to $\mathcal{R}=1$ can be enough to pass from an almost complete confinement, where most of the particles are advected in the tail, to an almost free escape. This means that energy independent arguments, in the description of how particles escape, are likely to be strongly inadequate, and only a detailed modeling of the full trajectories can give trustable results. \item Current sheets and lines are important confinement agents and they can strongly affect the pattern of escaping particles. This means that the structure, stability and dissipation of those features can substantially affect the way particles emerge out af these systems. In this respect the ability to properly model turbulence and dissipation, as key factors regulating those currents, becomes important for the level of anisotropy in the escape. \item Particle escape is likely to be charge separated (strong charge asymmetry) or partially charge separated. If the high energy particles that escape carry a sizable fraction of the overall energetics, this means that BSPWNe are likely to be characterized by the presence of charge separated flows. Charge separated flows, (and the related return currents) can give rise to filamentation, Weibel, and two-stream instabilities, which might affect deeply the underlying MHD structure that is based on charge neutrality. \end{itemize} One of the main issue related to the escape of high energy particle is that in systems like the Guitar nebula \citep{Hui_Becker07a} or the Lighthouse nebula \citep{Pavan_Bordas+14a} one sided bright X-ray feature are observed, which if interpreted in terms of high energy particles streaming out of the nebula into the ISM magnetic field \citep{Bandiera08a}, require a large level of anisotropy. Our results show that the required anisotropy can be due to the internal magnetic field. Unlike reconnection, which only affects low energy particles with Larmor radii smaller that the extent of the reconnection site, which in general are much smaller than the size of the bow-shock, global magnetic effects are relevant also for $\mathcal{R}\sim 1$, and they generate large anisotropy over the entire bow-shock head. Even of greater importance is the fact that the escaping particles might be in the form of a charge separated flow. This is highly relevant in terms of self confinement of the escaping particles. As the high energy particles start to stream in the ISM, they can drive streaming instability in a fashion not dissimilar to that of cosmic rays accelerated at SN shocks \citep{Ptuskin_Zirakashvili+08a,Malkov_Diamond+13a,Nava_Gabici+16a,DAngelo_Morlino+18a}. The level at which induced turbulence saturate depends strongly on the level of charge separation. In the presence of a net current, the magnetic field can be amplified much more efficiently that for a neutral flow \citep{Skilling71a,Bell04a}. This means that it is far easier to self confine the escaping particles in the vicinity of the BSPWN, if the outflow is charge separated. Indeed the recent TeV halo detected around Geminga \citep{Abeysekara_Albert+17a}, which has been interpreted as an evidence for a larger ISM turbulence that what is commonly assumed, could instead be due to local magnetic field amplification by a charge separated pair outflow \citep{Evoli_Linden+18a}. | 18 | 8 | 1808.08757 |
1808 | 1808.03390_arXiv.txt | Mergers of compact binaries (of a neutron star and another neutron star or a black hole, NSMs) are suggested to be the promising astrophysical site of the \rp. While the average coalescence timescale of NSMs appears to be $\gtrsim 100 \, \myr$, most of previous chemical evolution models indicate that the observed \SW{early appearance and large} dispersion of \rfe\ in Galactic halo stars at $\efeh \lesssim -2.5$ favors shorter coalescence times of 1--10 \myr. We argue that this is not the case for the models assuming the formation of the Galactic halo from clustering of subhalos with different star formation histories as suggested by Ishimaru et al. We present a stochastic chemical evolution model of the subhalos, in which the site of the \rp\ is assumed to be mainly NSMs with a coalescence timescale of $100 \, \myr$. \SW{In view of the scarcity of NSMs, their} occurrence in each subhalo is computed with a Monte Carlo method. Our results show that the less massive subhalos evolve at lower metallicities and generate highly \rp-enhanced stars. An assembly of these subhalos leaves behind the large star-to-star scatters of \rfe\ in the Galactic halo as observed. However, the observed scatters of [Sr/Ba] at low metallicities indicate the presence of an additional site that partially contributes to the enrichment of light neutron-capture elements such as Sr. The high enhancements of \rfe\ at low metallicities found in our low-mass subhalo models also \SW{qualitatively reproduce} the abundance signatures of the stars in the recently discovered ultra-faint dwarf galaxy Reticulum II. Therefore, our results suggest NSMs as the dominant sources of \rpes\ \SW{in the Galactic halo}. | Most of the elements with atomic numbers greater than $Z = 30$--40 are produced through neutron-capture processes, and about half of such heavy elements originate from the rapid neutron-capture process (\rp). However, the astrophysical site of the \rp\ has not been specified, which remains a long-standing problem in nuclear astrophysics (see, e.g., \citealt{Thielemann+17} for a recent review). Stellar abundances of Galactic halo stars serve as the fossils of the early history of the Galaxy, providing us with important clues to the Galactic formation and early chemical evolution. Metal-poor stars, those with metallicities of \feh% \footnote{$[{\rm A} / {\rm B}] \equiv \log_{10}(X_{\rm A} / X_{\rm B}) - \log_{10}(X_{\rm A, \odot} / X_{\rm B, \odot})$, where $X_{\rm A}$ and $X_{\rm B}$ are the mass fractions of elements $\rm A$ and $\rm B$.} $\lesssim -2.5$, are thought to be among the oldest objects in the Galaxy, which presumably have been formed in the first few hundred Myr of its life. Spectroscopic observations of metal-poor halo stars show a large star-to-star scatter of about 2.5~dex in the abundances of Eu (as representative of \rp\ elements) with respect to Fe, \eufe\ (e.g., \citealt{Honda+04, Francois+07, SAGA, Sneden+08}). In particular, several metal-poor stars such as CS~22892-052 \citep{Sneden+03} and CS~310812-001 \citep{Siqueira+13} show extremely high ratios of $\eeufe = 1.6$--1.7. These unique abundance signatures indicate that \SW{Eu does} not share common \SW{astrophysical conditions} with $\alpha$ or iron-group elements. It is also well known that such \rp-enhanced stars, which account for about 10\% of all metal-poor stars, exhibit fairly robust abundance distributions that agree with the solar system \rp\ pattern. On the one hand, the excellent agreement for the heavier side ($Z > 50$; \citealt{Sneden+08}) indicates the presence of the single robust ``main" \rp\footnote{Hereafter, we indicate the main \rp\ by the ``\rp", which produces all of the \rp\ elements with a solar \rp-like pattern but with a smaller content of $Z < 50$ elements.} site. On the other hand, the less remarkable agreement for the lighter side ($Z < 50$; \citealt{Siqueira+14}) as well as the bulk of (\rp-deficient) stars showing higher [Sr/Ba] than the solar \rp\ ratio \citep{McWilliam98, Johnson+02} implies the presence of another ``weak" \rp\ that produces only light neutron-capture elements \citep{Wanajo+06}. In fact, several metal-poor stars showing a descending trend of neutron-capture elements toward the heavier side have been identified \citep{Honda+06, Honda+07, Aokim+17}, which may reflect such a process. The major candidates suggested as the \rp\ site include core-collapse supernovae (CCSNe; e.g., \citealt{Burbidge+57, Hillebrandt+84, Woosley+94}) and binary mergers of a neutron star and another neutron star or a black hole (NSMs\footnote{In this paper, we generally indicate ``neutron star--neutron star mergers" by ``NSMs", although similar conclusions may be obtained for neutron star--black hole mergers.}; e.g., \citealt{Lattimer+74, Symbalisty+82, Eichler+89, Meyer89, Freiburghaus+99, Goriely+11, Korobkin+12, Wanajo+14}). Inhomogeneous chemical evolution models \citep{Ishimaru+99, Argast+00, Tsujimoto+00, Ishimaru+04, Cescutti08}, which account for incomplete mixing of nucleosynthetic ejecta in the interstellar medium (ISM), have demonstrated that the observed dispersion of \rfe\ ratios (where $r$ indicates an \rpe) in Galactic halo stars can be explained if CCSNe from a limited initial stellar mass range are the sources of \rp\ elements. However, recent nucleosynthesis studies show difficulties in producing heavy \rp \ elements $(Z \gtrsim 50)$ in the physical conditions relevant to CCSNe \citep{Wanajo+11, Wanajo+18, Wanajo13, Bliss+18}, which can be at best the sources of light neutron-capture elements made by a weak \rp. Effects of a strong magnetic field also have been discussed \citep{Thompson03, Suzuki+05, Winteler+12, Nishimura+15}, although their roles on the \rp\ are under debate \citep{Nishimura+17, Thompson+17, Moesta+17}. By contrast, recent nucleosynthesis calculations based on the hydrodynamical simulations of NSMs reasonably reproduce the solar $r$-process abundance curve (\citealt{Wanajo+14, Goriely+15, Radice+16, Wu+16}). The discovery of an electromagnetic emission \citep[kilonova;][]{Li+98, Metzger+10} associated with the gravitational-wave source GW170817 \citep{Abbott+17} also supports NSMs as the site of the \rp\ in the universe. In fact, the inferred amount of the \rp\ material ejected from this event, about 0.03--$0.05\, M_\odot$ \citep[e.g.,][]{Pian+17}, appears to be sufficient to account for the total mass of \rp\ elements in the Galaxy, provided that GW170817 is representative of NSM events. However, binary population synthesis models (e.g., \citealt{Dominik+12}) as well as observations of binary neutron stars \citep[e.g.,][]{Beniamini+16} suggest an average NSM coalescence timescale $\langle \etnsm \rangle \gtrsim 100 \, \myr$, which appears too long to allow for the observed appearance of, e.g., Eu at metallicities as low as $\efeh\sim -2.5$ on the basis of one-zone chemical evolution models \citep{Argast+04}. Such models suggest $\etnsm = 1-10 \, \myr$ to reproduce the observed early evolution of \rfe\ \citep{Argast+04, DeDonder+04, Matteucci+14, Tsujimoto+14, Cescutti+15, Wehmeyer+15}. Previous chemical evolution models based on the hierarchical structure formation of the Galactic halo also favor short NSM coalescence timescales (\citealt{Komiya+14, vandeVoort+15}). In addition, it is argued that the inferred low event rate ($0.4-77.4 \, \myr^{-1}$; \citealt{Dominik+12}) of NSMs causes too large \rfe\ dispersions at higher metallicity to be compatible with observations \citep{Qian00, Argast+04}. Chemical evolution studies assuming multiple $r$-process sites such as NSMs and magnetorotationally driven CCSNe attempt to explain the observed \rfe\ evolution of the metal-poor stars (\citealt{Wehmeyer+15, Shibagaki+16}). However, the uniqueness of the abundance patterns in \rp-enhanced stars apparently disfavors multiple \rp\ sites with different abundance distributions \SW{\citep[e.g.,][]{Sneden+08}} considered in these studies. \citet{Prantzos+06} suggested that the observed appearance of Eu at low metallicity as well as the large dispersion of [Eu/Fe] could be naturally explained if the Galactic halo was formed from merging subhalos with different star formation histories and if the production sites of Fe and \rp\ elements evolved on different timescales. \citet[][hereafter IWP15]{Ishimaru+15} have first explored this idea using a semianalytical model of merging subhalos, each of them evolving on a different timescale---depending on its mass---in a homogeneous way (i.e., the gas is assumed to be well mixed within each subhalo). According to IWP15, NSMs start occurring at the metallicity $\feh \lesssim -3$ and contributing to the enrichment of Eu even with the coalescence timescale of 100~Myr if less massive subhalos evolve with lower star formation efficiency. Similar results can be found in recent semianalytic \citep{Komiya+16} and chemodynamical \citep{Shen+15, Hirai+15, Hirai+17} studies \SW{(see also \citealt{Cote+17} for a comparison of several chemical evolution studies mentioned above with their own model)}. IWP15 also show that various star formation efficiencies make a difference in the numbers of cumulative NSMs occurring in subhalos. In fact, the number of cumulative average NSMs for the lightest subhalo with a stellar mass of $10^4\, M_\odot$ (similar to that of an ultra-faint dwarf galaxy, UFD) predicted in IWP15 is $\sim 0.1$, implying that only one 1 of 10 small subhalos experiences an NSM event. They suggested that a single NSM occurring in the least massive systems would lead to a very high \rfe\ of the inter stellar medium. The recently discovered UFD, Reticulum II \citep[Ret~II;][]{Ji+16_Nature, Roederer+16, Ji+16_Sr, Ji+18}, could be such an example as anticipated by IWP15, in which seven (out of nine) stars exhibited high \rfe\ ratios comparable to those in the most \rp-enhanced Galactic halo stars. In this paper, we extend the study of IWP15 to explain the presence of such \rp-enhanced stars and the scatter of \rfe\ ratios in the Galactic halo as well as in Ret~II. We also aim to examine the chemical evolution of Sr to test if our model is compatible with those of light neutron-capture elements. For this purpose, we construct a chemical evolution model, in which individual subhalos stochastically experience NSM events. In section \ref{sec:model}, the concept and setup of our model are presented in detail. In section \ref{sec:results}, we discuss the enrichment histories of Ba and Eu as representative of \rp\ elements. We also compare our results with the observations of a light neutron-capture element Sr (section~\ref{sec:strontium}) as well as of the \rp-enhanced stars in Ret II (section~\ref{sec:ufd}). Finally, we summarize and conclude our work in section \ref{sec:summary}. | \label{sec:summary} We have revisited the study of Galactic chemical evolution by IWP15 in order to investigate the role of NSMs as the dominant sources of \rp\ elements in the Galaxy. Our chemical evolution model was constructed on the basis of the hierarchical structure formation scenario in IWP15, in which the different mass ($M_\mathrm{sub}$) subhalos that formed the Galactic halo had different star formation histories. The number of NSMs occurring in each subhalo was obtained with the Monte Carlo method. The star formation histories of the subhalos were determined from the observed mass-metallicity relation of the local dwarf galaxies, assuming the same correlation for both systems. In the framework of our simple model of galactic chemical evolution, this mass-metallicity relation leads to $\ofr/\sfr \propto {\eMsub}^{-0.3}$, where $\ofr$ and \sfr\ are the outflow rate and the star formation rate, respectively. We examined two extreme cases as in IWP15, in which either of the coefficients for SFR and OFR (Equations~(1) and (2)) were kept constant such that $\ekof = 1.0 \, \gyr^{-1}$ and $\eksf \propto {\eMsub}^{+0.3}$ for case 1 or $\eksf = 0.20 \, \gyr^{-1}$ and $\ekof \propto {\eMsub}^{-0.3}$ for case 2. Our result shows that the observed properties of the enrichment histories of \rfe\ in the Galactic halo can be explained by assuming that NSMs are the sources of \rp\ elements. The adopted low SFR in low-mass subhalos for case~1 makes the occurrence of NSMs at low metallicities ($\feh\ \sim -3$) possible as shown by IWP15. In addition, the presence of the \rp-enhanced metal-poor stars in the Galactic halo is accounted for as a result of the large enhancement of \rfe\ by a single or a few NSMs in a small fraction of low-mass subhalos. However, our case 2 adopting a constant star formation efficiency ($k_\mathrm{SF}$) results in the enrichment of \rp\ elements at higher metallicities, $\feh\ > -2.5$, as found in previous studies. We conclude, therefore, that the reality is closer to our case~1, i.e., star formation is less efficient in lower-mass subhalos while gas outflow only weakly depends on the subhalo masses, provided that NSMs are the dominant contributors of \rp\ elements in the Galaxy. The observed trend of \rfe\ in the Galactic halo can be mostly reproduced solely by long-lived NSMs with a coalescence timescale of $100 \, \myr$. However, a small fraction of short-lived (1~Myr in our model) NSMs appear to be necessary, which are responsible for explaining the presence of stars with subsolar \rfe\ values at $\feh\ \lesssim -3$. A test shows that a long coalescence timescale appreciably greater than 100~Myr, such as $500 \, \myr$, has difficulty in reproducing the enhancement of \rpes\ at low metallicities. This implies either that in reality the distribution of coalescence timescales (for the long-lived NSMs) has a sharp peak at $\sim 100$~Myr or the NSMs with longer timescales escape from subhalos because of neutron star kicks and do not contribute to Galactic chemical evolution \citep{Beniamini+16b, Safarzadeh+17}. In the future, localizations of NSMs in galaxies by identification of electromagnetic counterparts (kilonovae) of gravitational waves will provide us with information on the distribution of coalescence timescales. It is important to note that our model naturally reproduces a large dispersion of abundance ratios (relative to Fe) for \rpes, but a small one for intermediate-mass elements such as Mg; the latter are produced by the same sources and on the same timescales as Fe, while the former result from different sources, operating on very different timescales. This feature, which is in good agreement with spectroscopic results, is another important aspect of NSMs as sources of \rpes. Our model is also successful in explaining the spectroscopic abundances of light neutron-capture elements such as $\sr$ when assuming an additional contribution (a weak \rp) from low-mass CCSNe. The resulting evolution of \srba\ as a function of metallicity reasonably reproduces the observed star-to-star scatters with higher [Sr/Ba] values than those predicted solely by the enrichment from NSMs. This supports the idea that there are extra sources (weak \rp) of light neutron-capture elements in addition to the (main) \rp. Finally, our models of the least massive subhalos with $M_\mathrm{sub} \sim 10^4\, M_\odot$ well account for the observed nature of UFDs, namely, only 1 (Ret~II) out of about 10 such galaxies shows enhancement of \rpes. Moreover, the fact that seven out of observed nine stars in Ret II show enrichment of \rpes\ can be reasonably reproduced by such a lowest-mass subhalo model. This supports the idea that the UFDs are the leftovers of the building blocks that made the Galaxy. Our models also predict the presence of Ret~II-like UFDs with various fractions of \rp-enhanced stars. This will be tested by future spectroscopic explorations of UFD galaxies. \SW{However, it is currently unclear if the mass-metallicity relation can be applied to the bulk of UFDs. It also is cautioned that we applied our homogeneous and continuous chemical evolution model to such a small system that probably had only a few episodes of star formation. Obviously, further studies of UFDs will be necessary from both observational and theoretical sides.} \SW{It should be noted that our study neglected a spatial inhomogeneity of ISM in each subhalo as well as a merging process of subhalos during the evolutionary time of the Galactic halo (2 Gyr). The contribution of $s$-process elements also was excluded, which could be important for the evolutions of Ba and Sr at [Fe/H] $\gtrsim -2.5$. Nevertheless, our simplified approach enabled us to disentangle the different sources of dispersion: the subhalo-mass-dependent SFRs that lead to a spread in [Fe/H] for the same [$r$/Fe] (as shown by IWP15) and the stochastic NSM events that lead to a spread in [$r$/Fe] for the same [Fe/H]. However, our models failed to explain the moderate dispersion of the observed [$r$/Fe] ratios at [Fe/H] $\sim -1.5$ (Figure~4), which may be due to the combined effects of ISM inhomogeneity and $s$-process contribution. Such effects will be explored in our forthcoming paper.} | 18 | 8 | 1808.03390 |
1808 | 1808.01489_arXiv.txt | \begin{flushleft} We analyze two different algorithms for constructing weakly inhomogeneous models for the low-redshift Universe, in order to provide a tool for testing the photon dynamics, within the sphere of validity for the Universe acceleration. We first implement the so-called quasi-isotropic solution in the late Universe, when a pure dark energy equation of state for the cosmological perfect fluid is considered. We demonstrate, that a solution exists only if the physical scale of the inhomogeneities is larger then the Hubble scale of the microphysics\footnote{Given a point in space $A$, the Hubble length may be considered as the radial distance of the points that, due to the Hubble expansion, recedes from $A$ with a speed equals to $c$, where $c$ is the speed of light.}, which implies that inhomogeneities could not be observed at present-stage time. Then, we analyze a weakly deformed isotropic Universe toward a spherically symmetric model, thought as the natural metric framework of the $\Lambda CDM$ model. The obtained picture offer a useful scenario to investigate the influence of the inhomogeneity spectrum (left free in the obtained solution), on the photon propagation at low redshift values.\\ \end{flushleft} \vspace{2mm} \emph{\bf Keywords} : Theoretical Cosmology, \and General Relativity, \and Numerical Cosmology, \and Low-redshift Universe, \and Weakly Inhomogeneous Universe, \and Low Redshift Universe | \label{intro} The Standard Cosmological model \cite{PrimCosm}\cite{KT90}, is based on the homogeneous and isotropic \emph{Robertson-Walker} metric, and bases its reliabillity on the high isotropy of the \emph{Cosmic Microwave Background Radiation} \cite{WMAP}\cite{Planck}.\\ Actually, the estimates from the galaxy surveys of the spatial scale at which the present Universe reaches homogeneity, provides a value of about $60 Mpc/h$ \cite{Cosmoscale}, where $h \approx 0.7$, which implies that, at lower scales, significant deviations from the Robertson-Walker geometry may be observed, at least as higher order corrections.\\ The presence of such small scale deviations, could influence the information we get from extragalactic sources and, in general, they affect the photon paths from distant regions up to our detectors.\\ In this work, we investigate the spatial metric which admit inhomogeneous corrections to the flat isotropic model (the present contribution of the spatial curvature is clearly negligible with respect to the matter terms).\\ As first step, we analyze the so-called quasi-isotropic solution \cite{LifKal} (see also \cite{SemiIsoMon}\cite{SIInflation}), but implemented in the low redshift Universe.\\ In particular, we study this solution in the presence of a perfect fluid, having a \emph{dark energy} \cite{darkenergy} equation of state, i.e. $P = w\rho$ ($P$ and $\rho$ being the pressure and energy density of the fluid respectively), where $-1 < w < -1/3$.\\ The interesting feature of the obtained solution, relies in the fact that the considered inhomogeneities, included in the model as small corrections, must correspond, in order the solution be consistent, to physical scales much greater than the \emph{Hubble (microphysical) scale}.\\ This fact makes such inhomogeneous corrections of pure curvature nature and, over all, they can not affect the physical processes taking place in the Hubble sphere.\\ Then, we consider the case of a \emph{Lemaitre-Tolmann-Bondi} spacetime \cite{PrimCosm}\cite{LTB}, which describes a spherically symmetric Universe in the presence of a matter source, and a non-zero \emph{cosmological constant}, an appropriate scenario to account for the so-called $\Lambda CDM$ model \cite{cosmocostant}.\\ Clearly, in order to describe the local behavior of the actual Universe, we consider the inhomogeneous perturbations again, as first order modification of the flat Robertson-Walker geometry.\\ We demonstrate the existence of a consistent solution of the linearized \emph{Einstein equations}, which does not fixes the radial dependence of the inhomogeneities, but only their time scaling.\\ Moreover, we specialize the obtained solution to the case of the $\Lambda CDM$ model, by tuning the values of the parameters in order to obtain, that the matter be the $30\%$ and the constant energy density the $70\%$ of the Universe critical parameter respectively.\\ The obtained time profile for the perturbations, together with the arbitrariness of their specific spatial morphology, offer an interesting arena to study the effects on the photon propagation, due to the local deviations of the actual Universe from homogenity.\\ Furthermore, the results of our analysis suggest an intriguing issue: while the inhomogeneities allowed by a $\Lambda CDM$ model are physically observable, living in principle, in the present Hubble sphere, the dark energy dominated Universe appears incompatible with the physical scale of inhomogeneity, the microphysics processes remain essentially concerned by the homogeneity restriction.\\ The different behavior of the two considered equations of state (we recall that the cosmological constant is associated to the relation $P = -\rho$), could become a qualitative discrimination property when incoming missions, like Euclid \cite{Euclid} will be able to test the large scale properties of the Universe, detecting details of the matter distribution across the cosmological space.\\ This paper is structured as follows : \\ on section \ref{sec:2.}, the \emph{Lifshitz-Khalatnikov} quasi isotropic solution, for a pure radiation high-redshift universe will be described, so that on section \ref{sec:3.}, the same procedure will be applied to the case of a low-redshift dark energy universe.\\ Following the same steps, on section \ref{sec:4.} is introduced the \emph{Lemaitre-Tolman-Bondi} model for spherically simmetric universes in the generic case, while in section \ref{sec:5.}, the previous solution will be extended to the case of a low redshift universe, filled with both a matter and a cosmological constant perfect fluid.\\ On section \ref{sec:6.}, the \emph{weakly inhomogeneous} model derived in the previous section, will be fitted with the actual observational data, in order to describe as best as possible the behaviour of our $\Lambda CDM$ universe.\\ Lastly, the article will be closed with concluding remarks, that are reported in section \ref{sec:7.}.\\ | \label{sec:8.} We analyze two different, but complementary algorithms to deal with small inhomogeneus corrections to the isotropic Universe: on one hand, we studied the so-called \emph{quasi-isotropic solution}, as implemented to a late dynamics, on the other one, we study a \emph{Lemaitre-Tolamann-Bondi} spherically symmetric solution, containing only small deviations depending on the pure radial coordinate.\\ We considered in both cases, sources in the form of a perfect fluid, but, while for the quasi-isotropic case we consider a dark energy equation of state with $-1< w<-1/3$, the spherically symmetric solution contains two different contribution, a matter fluid and a cosmological constant, respectively.\\ The basic result of our analysis, is demonstrating that the presence of a real dark energy contribution, prevent the possibility to deal with physical scales of the inhomogeneous correction being smaller than the actual Hubble scale of the Universe.\\ This constraint, comes from the necessity to rule out of the solution the spatial curvature contribution (due to inhomogeneous corrections), and it has very deep implications: the obtained perturbed solution is characterized by perturbations evolving only from a kinematical point of view, but unaffected by microphysical processes and, de facto they are not observable at the present time. Clearly, this result can not be considered as a general one for two basic reasons : \begin{itemize} \item It depends on the details of the solution construction i.e. retaining curvature effects of the perturbations more general regimes can exist; \item The perturbation component of a quasi-isotropic solution is naturally factorized into the product of a space dependent and a time dependent function respectively and this influences the obtained dynamics.\\ \end{itemize} For a discussion of related problems regarding sub-Hubble inhomogeneities see \cite{SubHub}.\\ The situation is different for the Lemaitre-Tolmann-Bondi model, where, we actually consider spherically symmetric deviations to the background dynamics, underlying the $\Lambda CDM$ model for the actual Universe.\\ We construct a inhomogeneous perturbation to the isotropic cosmology, whose spatial dependence, i.e. whose spectrum, is not fixed by the solution method, remaining a useful degree of freedom for fitting different physical situations.\\ Apart from the conceptual difference qualitatively emerging in the present study between the two used algorithms, the main merit of this work is outlining that in the \emph{LTB} case, correspond to a consistent solution with late time sub-Hubble inhomogeneities.\\ In fact, the possibility to check the photon dynamics on different weakly inhomogeneous Universe, offer an interesting tool to test some physical properties of the actual low redshift Universe.\\ In particular, as briefly shown in subsection \ref{sec:8.}, we suggest that the Lemaitre-Tolmann-Bondi model studied above, could be adopted to try to account with weak inhomogeneity profiles, the discrepancy existing between the value of the \emph{Hubble constant} $H_0$ \cite{Hubble}, as it is measure by \emph{WMAP} and \emph{Planck} Satellities and by the ground based surveys \cite{HDisc} \cite{Survey}.\\ The elimination of such a discrepancy by the proposed scenario, could put limit on the local inhomogeneity profile of the actual Universe, possibly tested by the incoming mission Euclid \cite{Euclid}. \\ | 18 | 8 | 1808.01489 |
1808 | 1808.08082_arXiv.txt | {We study the capture and subsequent annihilation of inelastic dark matter (DM) in the Sun, placing constraints on the DM-nucleon scattering cross section from the null result of the IceCube neutrino telescope.~We then compare such constraints with exclusion limits on the same cross section that we derive from XENON1T, PICO and CRESST results.~We calculate the cross section for inelastic DM-nucleon scattering within an extension of the effective theory of DM-nucleon interactions which applies to the case of inelastic DM models characterised by a mass splitting between the incoming and outgoing DM particle.~We find that for values of the mass splitting parameter larger than about 200 keV, neutrino telescopes place limits on the DM-nucleon scattering cross section which are stronger than the ones from current DM direct detection experiments.~The exact mass splitting value for which this occurs depends on whether DM thermalises in the Sun or not.~This result applies to all DM-nucleon interactions that generate DM-nucleus scattering cross sections which are independent of the nuclear spin, including the ``canonical'' spin-independent interaction.~We explicitly perform our calculations for a DM candidate with mass of 1 TeV, but our conclusions qualitatively also apply to different masses.~Furthermore, we find that exclusion limits from IceCube on the coupling constants of this family of spin-independent interactions are more stringent than the ones from a (hypothetical) reanalysis of XENON1T data based on an extended signal region in nuclear recoil energy.~Our results should be taken into account in global analyses of inelastic DM models.} \begin{document} | \label{sec:introduction} The existence of invisible mass, or dark matter (DM), in our Universe is supported by observations performed on very different physical scales.~These include the anomalous motion of stars and galaxies, gravitational lensing events in cluster of galaxies, patterns in the anisotropies of the cosmic microwave background radiation, and the observed large scale structure of the Universe (see~\cite{Bertone:2004pz} and references therein for a comprehensive review).~In the leading paradigm of modern cosmology DM is made of hypothetical particles with interactions at the weak scale or below~\cite{Arcadi:2017kky}.~This class of DM particles is currently searched for using, e.g., direct detection experiments~\cite{Goodman:1984dc}, which look for nuclear recoils induced by the non-relativistic scattering of DM particles in low-background detectors, and indirect detection experiments, which search for DM annihilation signals produced in space or at the centre of the Sun or of the Earth~\cite{Krauss:1985ks,Freese:1985qw}, where DM is expected to accumulate by losing energy while scattering off nuclei in the solar and terrestrial interiors.~Among the indirect detection experiments, neutrino telescopes, such as IceCube~\cite{Aartsen:2016zhm}, which search for neutrinos from DM annihilation in the Sun are of special interest for this work. Interpreting the null result of current experiments, DM is commonly assumed to scatter off nuclei elastically, i.e.~the DM particle is in the same state before and after scattering, e.g.~\cite{Lewin:1995rx}.~While this assumption is often fulfilled by popular models for DM~\cite{Bergstrom00}, it is not always true.~For example, in DM-nucleus collisions a DM particle could scatter to an excited state of higher mass (endothermic reaction), or scatter from an excited state to a different state of lower mass (exothermic reaction)~\cite{TuckerSmith:2001hy}\footnote{An alternative scenario is the one where the target nucleus is scattered off to an excited state~\cite{Goodman:1984dc}.~This scenario will not be discussed here.}.~The family of models where DM scatters off nuclei inelastically is collectively referred to as inelastic DM.~Inelastic DM has initially been proposed as an attempt to reconcile the observation of an annual modulation in the rate of nuclear recoil events recorded by the DAMA collaboration with the null result reported by other experiments~\cite{TuckerSmith:2001hy}.~In this context, it has also been noticed that inelastic DM-nucleus scattering can occur in a variety of theories, including supersymmetric models of nearly pure Higgsinos~\cite{Fox:2014moa}, magnetic inelastic DM~\cite{Chang:2010en}, and dark photon mediated DM~\cite{Smolinsky:2017fvb}.~While the initial motivation based on reconciling DAMA with the null result from other experiments has become less attractive due the strong exclusion limits presented by the LUX, XENON and PandaX collaborations, the fact that inelastic DM appears naturally in a variety of frameworks holds true.~Furthermore, it has been shown that within specific realisations of inelastic DM, the range of mass splittings between incoming and outgoing DM particles can be broader than initially proposed~\cite{Bramante:2016rdh}.~The large mass splitting limit of inelastic DM is known as the ``inelastic frontier''. The kinematics of inelastic DM-nucleus scattering is significantly different from the one of elastic interactions~\cite{Menon:2009qj,Bozorgnia:2013hsa,Scopel:2014kba,Scopel:2015eoh,Blennow:2015hzp,DelNobile:2015lxa}.~In particular, for DM particles heavier than atomic nuclei, e.g.~of mass 1 TeV, the inelastic DM-nucleus scattering is characterised by:~1) A finite minimum velocity the DM particle must have for the scattering to be kinematically allowed which scales like the inverse of the square root of the target nucleus mass;~2) A minimum nuclear recoil energy required for the scattering to occur which is approximately equal to the mass splitting between the incoming and outgoing DM particle.~These properties imply that inelastic DM-nucleus scattering is kinematically favoured for target nuclei with large mass numbers, and that only direct detection experiments which record nuclear recoil energies larger than the inelastic DM mass splitting parameter can effectively probe this scenario.~Based on these properties, it has been found in~\cite{Bramante:2016rdh} that an experiment like CRESST, which probes a range of nuclear recoil energies larger than, e.g., XENON1T, is effective in setting limits on the DM-nucleon scattering cross-section in the large mass splitting limit.~On the other hand, as far as neutrino telescopes are concerned, the inelastic frontier of DM models remains as of yet unexplored. In this article we set constraints on the DM-nucleon scattering cross section of inelastic DM models in the large mass splitting limit using data from neutrino telescopes~\cite{Aartsen:2016zhm} and direct detection experiments~\cite{Aprile:2018dbl,Amole:2015pla,Angloher:2015ewa}.~The DM-nucleon scattering cross section is computed within the non-relativistic effective theory of DM-nucleon interactions, formulated in~\cite{Fitzpatrick:2012ix}, applied to the analysis of neutrino telescope data and DM capture in the Sun and Earth in~\cite{Liang:2013dsa,Catena:2015uha,Catena:2015iea,Catena:2016ckl,Catena:2016kro,Kavanagh:2016pyr,Widmark:2017yvd}, and extended to inelastic DM in~\cite{Barello:2014uda}.~Our constraints from neutrino telescopes are compared with those we obtain from an analysis of XENON1T, PICO and CRESST results.~We find that in the inelastic frontier, exclusion limits from neutrino telescopes can be stronger than those from direct detection, even for canonical spin-independent DM-nucleon interactions.~This result should be taken into account in the analysis of IceCube data within inelastic DM models. This article is organised as follows.~In Sec.~\ref{sec:inelastic} we introduce the theory of inelastic DM, focusing on kinematical aspects (Sec.~\ref{sec:kinematics}), and on the expected signals at direct detection experiments (Sec.~\ref{sec:dd}) and neutrino telescopes (Sec.~\ref{sec:nt}).~Sec.~\ref{sec:results} focuses on our limits from neutrino telescopes and direct detection experiments on the DM-nucleon scattering cross section of inelastic DM models in the large mass splitting limit.~We comment on our results and conclude in Sec.~\ref{sec:conclusion}. | \label{sec:conclusion} We have studied the capture and subsequent annihilation of inelastic DM in the Sun, placing constraints on the DM-nucleon scattering cross section (which is quadratic in $c_j^0$ and $c_j^1$) from the null result of IceCube.~The cross section for inelastic DM-nucleon scattering has been calculated within an extension of the effective theory of DM-nucleon interactions which applies to the case of inelastic DM.~We have explicitly performed our calculations assuming a DM particle mass of 1 TeV, but our conclusions qualitatively also apply to DM particle candidates with different masses.~We find that for values of the mass splitting parameter larger than about 200 keV neutrino telescopes place limits on the DM-nucleon scattering cross section which are stronger than the ones from current DM direct detection experiments.~The exact mass splitting value depends on whether DM thermalises in the Sun or not.~This result applies to all DM-nucleon interactions that generate DM-nucleus scattering cross sections which do not depend on the nuclear spin, including the ``canonical'' spin-independent interaction, i.e.~operator $\hat{\mathcal{O}}_1$ in Tab.~\ref{tab:operators}.~Indeed, for these interactions IceCube exclusion limits on the corresponding coupling constants remain relatively flat up to mass splittings of about 300 keV.~Furthermore, we find that exclusion limits from IceCube on the coupling constants of this family of interactions are more stringent than the ones from a (hypothetical) reanalysis of XENON1T data based on an extended signal region in nuclear recoil energy.~Our results should be taken into account in the analysis of neutrino telescope data, and in global statistical analysis of inelastic DM models. | 18 | 8 | 1808.08082 |
1808 | 1808.06867_arXiv.txt | {The Wilson depression is the difference in geometric height of unit continuum optical depth between the sunspot umbra and the quiet Sun. Measuring the Wilson depression is important for understanding the geometry of sunspots. Current methods suffer from systematic effects or need to make assumptions on the geometry of the magnetic field. This leads to large systematic uncertainties of the derived Wilson depressions.} {We aim at developing a robust method for deriving the Wilson depression that only requires the information about the magnetic field that is accessible from spectropolarimetry, and that does not rely on assumptions on the geometry of sunspots or on their magnetic field.} {Our method is based on minimizing the divergence of the magnetic field vector derived from spectropolarimetric observations. We focus on large spatial scales only in order to reduce the number of free parameters.} {We test the performance of our method using synthetic Hinode data derived from two sunspot simulations. We find that the maximum and the umbral averaged Wilson depression for both spots determined with our method typically lies within 100~km of the true value obtained from the simulations. In addition, we apply the method to Hinode observations of a sunspot. The derived Wilson depression ($\sim 600$~km) is consistent with results typically obtained from the Wilson effect. We also find that the Wilson depression obtained from using horizontal force balance gives 110 \-- 180~km smaller Wilson depressions than both, what we find and what we deduce directly from the simulations. This suggests that the magnetic pressure and the magnetic curvature force contribute to the Wilson depression by a similar amount.} {} | The geometric height at which unity continuum optical depth is reached is depressed within sunspots relative to the quiet Sun. This so-called Wilson depression~\citep{1774RSPT...64....1W} is caused by a lower opacity within the sunspot due to the lower temperature and a reduced gas pressure. The magnetic pressure and the curvature force of the strong magnetic field of sunspots balance this reduced gas pressure with the gas pressure of the surrounding quiet Sun. The connection between the Wilson depression and the strength and geometry of the magnetic field makes the Wilson depression an important quantity for understanding the structure of sunspots. In particular, it is not known by how much the curvature force of the magnetic field contributes to stabilizing the sunspot. Unfortunately, the Wilson depression remains one of the more poorly known parameters of sunspots. Several studies have tried inferring the Wilson depression by making use of the horizontal force balance between the sunspot and the surrounding quiet Sun. However, since it is unknown by how much the curvature force contributes to the force balance, an accurate estimate of the Wilson depression with this method is not possible. Depending on the assumed influence of the curvature force, the derived Wilson depression lies in the range between 400-1000~km \citep{1993A&A...277..639S,1993A&A...270..494M,2004A&A...422..693M}. The Wilson depression can also be estimated geometrically when the sunspot approaches the limb (i. e. via the Wilson effect). However, this method is influenced by radiative transfer or changes of the size of the umbra and penumbra with height \citep[see, e.g., the discussion in][]{2003A&ARv..11..153S}. Hence, the Wilson depression cannot be inferred very accurately when using this method. \citet{1972SoPh...26...52G} derived an average Wilson depression $z_{\rm W}$ of $600\pm 200$~km based on the Wilson effect. In contrast, \citet{1974SoPh...35..105P} measured a significantly larger Wilson depression of $950\--1250$~km with large spots having a higher Wilson depression ($z_{\rm W} = 1500\--2100$~km) than small spots ($z_{\rm W} = 700\--1000$~km). However, later results ($z_{\rm W} = 500 \-- 1000$~km) obtained by \citet{1983SoPh...88...71B} are in better agreement with those of \citet{1972SoPh...26...52G}. Here we present an alternative method for measuring the Wilson depression that does not need to make any assumptions on the geometry of the magnetic field in sunspots or on the structure of the sunspot. Our method is based on imposing the divergence-free condition on the magnetic field vector deduced from the inversion of observed Stokes profiles. This approach has already been used by \cite{2010ApJ...720.1417P} to derive the small-scale corrugation of the $\tau=1$ layer of a small patch within the penumbra of a sunspot. Here, we modify this method to provide the large-scale corrugation of the $\tau=1$ layer within the entire sunspot. We first perform a test of the method: we use it to derive the Wilson depression from synthetic Hinode observations generated from two MHD simulations of sunspots \citep{2012ApJ...750...62R,2015ApJ...814..125R}. After that we apply it to one spot observed with Hinode. We also compare our results with the Wilson depression derived using the pressure method. | We could successfully reproduce the Wilson depression of two simulated sunspots by minimizing the divergence of the magnetic field. We also applied this method to a sunspot observed with Hinode (AR~10923) and derived a Wilson depression that is consistent with the results statistically provided by some studies of the Wilson effect \citep[$\sim 600$~km,][]{1972SoPh...26...52G}. Our results suggest that the pressure due to the vertical component of the magnetic field and the curvature integral contribute by a similar amount to the horizontal force balance within sunspots, as has already been suggested by previous investigations \citep{1993A&A...270..494M,1993A&A...277..639S,2004A&A...422..693M}. The Wilson depression of the sunspot simulations by \citet{2012ApJ...750...62R,2015ApJ...814..125R} is smaller by $\sim 50$~km than the one of AR~10923. In the simulated spots, the curvature integral is significantly lower than the pressure due to the vertical component of the magnetic field, leading to a lower Wilson depression. This might be caused by the different geometry of the simulated spots compared to AR~10923 (see Table~\ref{tab:param}). The simulated spots exhibit a slightly stronger magnetic field but they are significantly smaller than the AR~10923 spot. The main limitation of our method is that it requires a reliable inversion of the full magnetic field vector over a broad range in height. In the umbra, our method requires an estimate of the magnetic field up to $\sim 800$~km above the $\tau = 1$ surface. The lines used in this study are not very sensitive at these height and so, the magnetic field is predominantly extrapolated from lower layers. Observations of multiple spectral lines with a broad range of formation heights are needed in order to retrieve a better inversion (as is planned, e. g., with the upcoming Sunrise III mission). Systematic errors also arise from the assumption of LTE and hydrostatic equilibrium in the inversion, although these are likely small compared with the uncertainty resulting from the inversions. Also, modeling the height dependency of atmospheric parameters with splines is not always a good representation of the true atmosphere, especially when there are strong gradients with height. In addition, the $180^\circ$-ambiguity of the Zeemann-effect needs to be resolved accurately across the entire spot at all optical depths. Fortunately, some of these inaccuracies of the inversion occur on small spatial scales, so that their influence on large spatial scales should be limited. For example, as shown in Section~\ref{sect:method}, neglecting the small-scale corrugations of the $\tau=1$ layer has only a small influence on the large-scale divergence of the magnetic field. A detailed determination of the errors in the spatially coupled inversion is beyond the scope of this paper. We therefore use the tests made with the synthetic data as a rough guide to the total error. The Wilson depression derived from the synthetic Hinode data exhibits an error of the order of $\sim 95$~km. In case of real observations, the error is probably somewhat higher. The synthetic data were generated using the same radiative transfer code and the same line parameters as were used for the inversion. Hence, the inversion of the synthetic data is likely to be more accurate than for real observations. Bearing these arguments in mind, we assign a preliminary error of 100~km to the Wilson depression derived from the Hinode observations. The estimated error of our method is mainly statistical in nature. Any systematic component is significantly lower than the one of the Wilson effect or of the pressure method. As explained in the introduction, the values of the Wilson depression derived using the Wilson effect vary by more than 1000~km between different studies in a systematic manner. In case of the pressure method, the inferred Wilson depression depends on the assumed value of the curvature integral. Our method is limited to the large-scale corrugations of the $\tau = 1$ surface. Measuring the Wilson depression with a higher spatial resolution requires including more Fourier coefficients in the Fourier series of the Wilson depression, which affects the minimization of the merit function. In addition, the divergence of the magnetic field is dominated by the largest spatial scales. When evaluating the divergence of the magnetic field for AR~10923 on a corrugated grid in geometrical height (300~km above $\tau = 1$ at each pixel), 76\% of the total variance of the divergence of the magnetic field occur within the range of spatial scales that we considered in Section~\ref{sect:AR10923} ($k_{\rm max} = l_{\rm max} = 3$). Hence, the merit function is not very sensitive to smaller spatial scales, at least when considering entire sunspots. | 18 | 8 | 1808.06867 |
1808 | 1808.01730_arXiv.txt | Early-type galaxies (ETGs) are composed of two distinct populations: high-mass and low-mass, which are likely to be built via gas-poor merging and gas-rich merging/accretion, respectively. However, it is difficult to directly associate low-mass ETGs with gas-rich processes, because currently they are gas poor with no signs of ongoing star formation. We report a discovery of an ETG (SDSS J142055.01+400715.7) with $M_{*}$=10$^{10}$ M$_{\odot}$ that offers direct evidence for gas-rich merging as the origin of low-mass ETGs. The integrated properties of the galaxy are consistent with a typical low-mass ETG, but the outer and inner regions show distinct dispersion- and rotation-dominated kinematics, respectively. There are some tidal features surrounding the galaxy. These two facts suggest very recent galaxy merging. Furthermore, the inner disk harbors on-going star formation, indicating the merging to be gas rich. This type of galaxy is rare but it may be a demonstration of the role the transient phase of gas-rich merging plays in making a low-mass ETG. | The Hubble sequence of galaxies as first proposed by \cite{Hub:26,Hub:36} divides galaxies into two types: late-type and early-type galaxies (ETGs), characterised by whether or not they contain spiral features. ETGs are traditionally classified into ellipticals and S0s depending on the presence of a disk, but the photometric analysis of galaxies' isophotes suggests that some ellipticals can still contain weak disk components \citep{Ben:89}, while inclination effects can lead to the mis-classification of some S0s as ellipticals \citep{Jor:94}. Among ETGs, high-mass ($M_{*}$ $>$ 10$^{11}$ M$_{\odot}$, $M_{\rm V}$ $\lesssim$ -21.5) and low-mass galaxies are found to be two distinct galaxy populations. Massive ETGs always have boxy and triaxial isophotal shapes while low-mass ones are disky \citep{Kor:99, Kor:09} and axisymmetric \citep{Cap:13a}. In terms of the radial profile of the optical light, massive ETGs have relatively large S\'ersic indices at the outer radius \citep[$\rm n\geq4$][]{Kra:13a} but show a deficit at the inner radius, known as the core. On the other hand, low-mass ETGs show relatively small S\'ersic indices at the outer radius ($\rm n\leq3$ \citet{Kra:13a}), and are core-less with either a power-law inner profile or even extra nuclear light \citep{Kor:99}. Observations further indicate that some central components have younger stellar populations than the rest of the galaxies \citep{Lau:05, Mcd:06}. With integral field unit (IFU) spectroscopic observations such as the SAURON survey \citep{deZ:02} and ATLAS$^{\rm 3D}$ project \citep{Cap:11a}, it is possible to classify ETGs according to their 2D kinematic properties. For example, \cite{Kra:06,Kra:11} found that ETGs can be separated into ``regular rotators'' with a regular velocity field and ``non-regular rotators'' dominated by random motion, associated with high-mass and low-mass ETGs, respectively. \citet{Ems:07} used the apparent angular momentum parameter to separate ETGs and found that massive ETGs rotate slowly (slow rotators) and low-mass ellipticals rotate faster (fast rotators). \cite{Cap:11a,Cap:11b} argued that fast rotators form a smooth parallel sequence to spiral galaxies on the luminosity/mass vs. size plane. It is suspected that the cores of massive ETGs form through dissipationless dry mergers in which supermassive black holes sink into the galaxy center to eject stars and form a nuclear core \citep{Kor:09, Cap:16, Kra:13b}, while extra nuclear light in low-mass ETGs is produced by gas-rich merging or accretion, in which a nuclear starburst is triggered. However, all ETGs, remnants of both gas-poor and gas-rich merging, lack on-going star formation since even in the case of gas-rich merging the gas has been dispelled or consumed, ceasing star formation. This makes it difficult to directly associate a low-mass ETG with gas-rich processes. In the MaNGA project (Mapping Nearby Galaxies at Apache Point Observatory ), we identified an ETG SDSS J142055.01+400715.7 ({\it RA=14h20m55.014s \& DEC=+40d07m15.70s}), which may be a living example that gas-rich merging is playing a major role in forming a low-mass ETG with extra nuclear light. In section 2 we briefly introduce the MaNGA project, the characteristics of the galaxy we focus on and the methodology. In section 3 we list the physical properties of this galaxy. Followed the discussion in section 4, the summary and conclusion is in section 5. | \label{sect:conclu} We report a discovery of an early-type galaxy in MaNGA that offers direct evidence for gas-rich merging as the origin of low-mass ETGs. The integrated properties of the galaxy are consistent with a typical low-mass ETG with $M_{*}$=10$^{10}$ M$_{\odot}$, no spiral arms, fast rotating kinematics and an outer S\'ersic profile index of 2.1-2.4. However, the outer and inner parts show distinct dispersion dominated and rotation dominated kinematics, respectively. And the broad-band image shows some tidal features surrounding the galaxy. These all suggest recent galaxy merging. Furthermore the inner disk harbors on-going star formation as indicated by the BPT diagram, which suggests that the merging is gas rich. The type of the galaxy is rare but it may be a demonstration of the role that gas-rich merging plays in making a low-mass ETG.\\ {\noindent \bf Acknowledgements} We thank the referee for a detailed report which helped significantly improve the presentation of our work. S.L. and Y.S. acknowledge support from the National Key R\&D Program of China (No. 2018YFA0404502), the National Natural Science Foundation of China (NSFC grants 11733002 and 11773013), the Excellent Youth Foundation of the Jiangsu Scientific Committee (BK20150014), and National Key R\&D Program of China (No. 2017YFA0402704). Y.C. acknowledges support from National Natural Science Foundation of China (NSFC grants 11573013). D.B. is supported by grant RScF 14-50-00043. Funding for the Sloan Digital Sky Survey IV has been provided by the Alfred P. Sloan Foundation, the U.S. Department of Energy Office of Science, and the Participating Institutions. SDSS- IV acknowledges support and resources from the Center for High-Performance Computing at the University of Utah. The SDSS web site is www.sdss.org. SDSS-IV is managed by the Astrophysical Research Consortium for the Participating Institutions of the SDSS Collaboration including the Brazilian Participation Group, the Carnegie Institution for Science, Carnegie Mellon University, the Chilean Participation Group, the French Participation Group, Harvard-Smithsonian Center for Astrophysics, Instituto de Astrof\'{i}sica de Canarias, The Johns Hopkins University, Kavli Institute for the Physics and Mathematics of the Universe (IPMU) / University of Tokyo, Lawrence Berkeley National Laboratory, Leibniz Institut f\"{u}r Astrophysik Potsdam (AIP), Max-Planck-Institut f\"{u}r Astronomie (MPIA Heidelberg), Max-Planck-Institut f\"{u}r Astrophysik (MPA Garching), Max-Planck-Institut f\"{u}r Extraterrestrische Physik (MPE), National Astronomical Observatory of China, New Mexico State University, New York University, University of Notre Dame, Observat\'{o}rio Nacional / MCTI, The Ohio State University, Pennsylvania State University, Shanghai Astronomical Observatory, United Kingdom Participation Group, Universidad Nacional Aut\'{o}noma de M\'{e}xico, University of Arizona, University of Colorado Boulder, University of Oxford, University of Portsmouth, University of Utah, University of Virginia, University of Washington, University of Wisconsin, Vanderbilt University, and Yale University. | 18 | 8 | 1808.01730 |
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