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1808 | 1808.06821_arXiv.txt | A flare kernel associated with a C4 class flare was observed in a spectral window including the He \emissiontype{I} triplet 10830 \AA\, and Si \emissiontype{I} 10827\AA\, with a spectropolarimeter on the Domeless Solar Telescope at Hida Observatory on August 9th, 2015. Observed Stokes profiles of the He \emissiontype{I} triplet in the flare kernel in its post-maximum phase are well reproduced through inversions considering the Zeeman and the Paschen-Back effects with a three slab model of the flare kernel, in which two slabs having up and downward velocities produce emissions and one slab produces an absorption. The magnetic field strength inferred from the emission components of the He \emissiontype{I} line is 1400 G, which is significantly stronger than 690 G that is observed at the same location in the same line 6.5 hours before the flare. In addition, photospheric magnetic field vector derived from the Si \emissiontype{I}10827\AA\,is similar to that of the flare kernel. To explain this results, we suggest that the emission in the He \emissiontype{I} triplet during the flare is produced in the deep layer, around which bombardment of non-thermal electrons leads to the formation of a coronal temperature plasma. Assuming a hydrogen column density at the location where the He \emissiontype{I} emissions are formed, and a power-law index of non-thermal electron energy distribution, we derived the low-energy cutoff of the non-thermal electron as 20 - 30 keV, which is consistent with that inferred from hard X-ray data obtained by RHESSI. | \label{sec.intro} Solar flares are sudden brightenings in the entire electromagnetic solar spectrum, from radio to $\gamma$-rays \citep{fletcher11}. Their phenomenological aspects are explained as the release of magnetic energy stored in the solar corona through magnetic reconnections (\cite{carmichael64}, \cite{sturrock66}, \cite{hirayama74}, \cite{kopp76}, \cite{nishio97}, \cite{melrose97}, \cite{shibata99}). Evaluation of amount of the magnetic energy that is converted to other forms of energy, such as non-thermal electrons, heat, and plasma motions, is an important subject for understanding the mechanism of the energy release \citep{emslie12, aschwanden17}. Solar chromosphere dynamically responds to solar flares through the energy injection from the corona by thermal conduction and/or non-thermal particles \citep{nagai80, somov81, somov82, ichimoto84}. For the thermal conduction, the transition region between the corona and the chromosphere will propagate downward due to the heating and subsequent evaporation of the upper part of the chromosphere. For the non-thermal particles, temperature may increase at a location in the chromosphere at which the non-thermal electrons penetrate and dissipate large part of their kinetic energy \citep{nagai84}. Therefore, the dynamical responses of the chromosphere to the flare will be different against the different heating mechanisms, and the investigation of it will provide us crucial information about the mechanism of energy transport in flares and their energetics. The polarization of photospheric or chromospheric spectral lines is produced by magnetic fields through the Zeeman and Hanle effects. Furthermore, when energetic particles with anisotropic velocity distribution, e.g., the accelerated charged particles and heat flux, penetrate into the chromosphere, collisions with ambient atoms will produce the impact polarization depending on the kinetic energy and the velocity distribution of the particles \citep{henoux83a, henoux83b}. Radiation anisotropy due to plasma inhomogeneities at the boundary of the flare ribbons also can induce scattering polarization in the spectral lines \citep{stepan13}. A number of works have reported observations of the polarization in spectral lines associated with flares, and they have discussed the origin of the polarization \citep{henoux83a, hanaoka03, xu05, judge15}. According to the current flare models (see e.g. \cite{shibata11}), flare ribbons observed in the chromosphere correspond to foot points of the reconnecting and restructuring coronal magnetic loops. Hence, measurements of the magnetic field strength in the flare ribbons allow us to estimate the reconnecting magnetic flux, and the released magnetic energy \citep{isobe02, asai04}. In addition, inferred magnetic field variations in photosphere during flares \citep{wang92, sudol05} are interpreted as a result of coronal field restructuring \citep{anwar93, hudson08}. Through spectro-polarimetric observations of chromospheric line, \citet{penn95} and \citet{harvey12} derived magnetic field component along the line of sight from circular polarization, i.e. Stokes ${\it V}$, of the He \emissiontype{I} 10830 \AA\,and Ca \emissiontype{II} 8542 \AA\,, respectively. \citet{kleint17} analyzed Stokes ${\it V}$ imaging spectro-polarimetric data in the Ca \emissiontype{II} 8542 \AA\, and found that a significant decrease and increase of chromospheric magnetic field component along the line of sight during a flare, and its variation is twice as large as that in the underlying photosphere. They attributed the change of the chromospheric magnetic field to the coronal field restructuring during the flare. However, since the direction of the magnetic fields was not measured, it remains unresolved if the coronal field restructuring was the major cause of the large changes in the observed fields. \citet{kuridze18} applied spectro-polarimetric inversions \citep{socas-navarro15} to Ca \emissiontype{II} 8542 \AA\,full-Stokes profiles of a flare ribbon, and suggested that the line in the flaring atmosphere is formed in the deeper layers than in the quiet Sun. They also present strong $\sim 1000$ - $1500$ G transverse component of the magnetic field in a portion of the flare ribbon where linear polarization is the strongest. This paper presents the first measurement of magnetic field vectors, i.e., directions and strengths, in a flare kernel with a spectro-polarimetric observation of the He \emissiontype{I} 10830 \AA. In the following sections, we describe the details of the observation (section \ref{sec.observation}), the results on the magnetic field vectors derived from Stokes profiles of the He \emissiontype{I} 10830 \AA\,in the flare kernel (section \ref{sec.result}), and finally we discuss possible interpretations and summarize our findings (section \ref{sec.sammarry}). | \label{sec.sammarry} We present the first measurement of the magnetic field vectors on the flare kernel with a spectro-polarimetric observation of the He \emissiontype{I} 10830 \AA. The observed Stokes profiles of the flare kernel in the He \emissiontype{I} 10830 \AA\,are well reproduced by the triple slab model that consists of two emission components with an upward and downward Doppler velocities and one absorption component. Another important findings are that the magnetic field strengths of the emission components are two times larger than that inferred from the Stokes profiles in the He \emissiontype{I} 10830 \AA\, before the flare occurrence, and four times larger than that of the absorption component. One of possible interpretations for this result is that the emissions of the He \emissiontype{I} 10830 \AA\,originate from the heating in the deep layer of the flare kernel due to bombardment of non-thermal electrons (figure \ref{fig.interpret}). When non-thermal electrons penetrates the chromosphere, they will stop at a depth determined by their kinetic energy and the ambient density, and it will lead to the formation of a coronal temperature plasma in the energy-deposited layer \citep{nagai84, tei18}. As a result, the He \emissiontype{I} 10830 \AA\,will be produced either ways by direct excitation by the non-thermal electrons in the chromosphere, collisional excitation by thermal electrons in compressed chromosphere next to the coronal temperature plasma, or the photoionization by EUV radiation with a wavelength $< 504$ \AA\, emitted from the coronal temperature plasma \citep{ding05, zeng14}.% The source of the observed two emission components with an upward and downward velocity (the fourth and fifth rows of the table \ref{tbl.result1}) can be attributed to the top and bottom boundary layers of the coronal temperature plasma (L2 and L3 in the Figure \ref{fig.interpret}), and the source of the absorption component (the third row of the table \ref{tbl.result1}) can be in the upper chromosphere, through which the non-thermal electron beam propagates without dissipation (L1 in the figure \ref{fig.interpret}). \begin{figure} \begin{center} \includegraphics[width=14cm]{interpret_07.eps} \end{center} \caption{ An interpretation of the inferred magnetic field vectors on the flare kernel. The dotted lines indicate the line of sight. } \label{fig.interpret} \end{figure} The magnetic field strength may decrease with the height in the flare kernel, because the field strength in the upper photosphere (the sixth row of the table \ref{tbl.result1}) is four times larger than that in the upper chromosphere (the third row of the table \ref{tbl.result1}). The magnetic field strengths of the observed two emission components of the He \emissiontype{I} 10830 \AA\, are also much larger than that in the upper chromosphere during and 6.5 hours before the flare (compare the first, fourth, fifth, and sixth rows of the table \ref{tbl.result1}). It is consistent with the scenario in which the emissions are produced in the deep layer. The depth where the non-thermal electrons deposit the most energy depends on the low-energy cutoff in their energy spectrum \citep{emslie78}.% We can estimate the low-energy cutoff by assuming the hydrogen column density in the pre-flare chromosphere \citep{allred15}. In non-flaring regions, the He \emissiontype{I} 10830 \AA\, line is supposed to be formed in the upper chromosphere \citep{avrett94}, and the difference of the formation heights of the He \emissiontype{I} 10830 \AA\,and the Si \emissiontype{I} 10827 \AA\,lines are estimated to be approximately 1000 km in sunspot umbrae from the analysis of propagating waves \citep{centeno09}. If we assume that the non-thermal electron beam with the power-law index of 7.8 (section \ref{sec.observation}) penetrate into the chromosphere and stop at the depth with a hydrogen column density of $1$ - $3 \times 10^{20}\,{\rm cm^{-2}}$, which corresponds to the depth of 1000 km below the transition region according to the umbral model of \citet{maltby86}, the low-energy cutoff can be derived as $\sim$20-30 keV (see figures 5, 6, and 9 of \cite{allred15}). Because the Si \emissiontype{I} 10827 \AA\,did not show emission, the hydrogen column density of $1$ - $3 \times 10^{20}\,{\rm cm^{-2}}$, or the derived lower cutoff energy, will be the maximum possible value. The hard X-ray spectral fit yields that the upper limit of the low-energy cutoff is $21.7 \pm 1.4$ keV (section \ref{sec.observation}), and it is consistent with the result here. It must be noted that there may be other possible scenarios, such as the coronal field restructuring and change of line formation height due to thermal conduction. In these scenarios, three components in the triple slab model can be interpreted as three different unresolved magnetic loops. The observed Stokes-${\it Q}$ and ${\it U}$ profiles are qualitatively different from those reported in a X-class flare by \citet{judge15}. They concluded that the linear polarization in He \emissiontype{I} 10830 \AA\,in the X-class flare are naturally explained with the scattering polarization under an anisotropic radiation field, while we successfully reproduced our observed profiles by the Zeeman effect without taking the scattering polarization into account. It might be possible that the kernels observed in these two flares are located in different conditions of the radiation fields and the magnetic fields at the He \emissiontype{I} 10830 \AA\, formation layer. Finally we like to stress that detailed measurements of temporal and spacial distributions of magnetic fields and plasma dynamics in flare kernels have a potential to provide valuable diagnoses of the non-thermal electrons in flares, and such observations at much higher spatial resolutions are expected to be performed by new generation instruments, such as Daniel K. Inouye Solar Telescope (DKIST, \cite{rimmele05}). \bigskip \begin{ack} This work was supported by JSPS KAKENHI Grant Number 22244013, JP16H01177, JP15K17609, 26560079, and 16H03955. We thank staffs and students of Ibaraki University and Kyoto University, especially Dr. A. Asai, Dr. S. Nagata, Dr. T. T. Ishii, Ms. A. Tei, and Mr. T. Nakamura. \end{ack} | 18 | 8 | 1808.06821 |
1808 | 1808.04223_arXiv.txt | {Orbital mean motion resonances in planetary systems originate from dissipative processes in disk-planet interactions that lead to orbital migration. In multi-planet systems that host giant planets, the perturbation of the protoplanetary disk strongly affects the migration of companion planets.} {By studying the well-characterized resonant planetary system around GJ 876 we aim to explore which effects shape disk-driven migration in such a multi-planet system to form resonant chains.} {We modelled the orbital migration of three planets embedded in a protoplanetary disk using two-dimensional locally isothermal hydrodynamical simulations. In order to explore the effect of several disk characteristics we performed a parameter study by varying the disk thickness, $\alpha$ viscosity, mass as well as the initial position of the planets. Moreover, we have carefully analysed and compared simulations with various boundary conditions at the disk's inner rim.} {We find that due to the high masses of the giant planets in this system, substantial eccentricity can be excited in the disk. This results in large variations of the torque acting on the outer lower-mass planet, which we attribute to a shift of Lindblad and corotation resonances as it approaches the eccentric gap that the giants create. Depending on disk parameters, the migration of the outer planet can be stopped at the gap edge in a non-resonant state. In other models, the outer planet is able to open a partial gap and to circularize the disk again, later entering a 2:1 resonance with the most massive planet in the system to complete the observed 4:2:1 Laplace resonance.} {Disk-mediated interactions between planets due to spiral waves and excitation of disk eccentricity by massive planets cause deviations from smooth inward migration of exterior lower mass planets. Self-consistent modelling of the disk-driven migration of multi-planet systems is thus mandatory. Constraints can be placed on the properties of the disk during the migration phase, based on the observed resonant state of the system. Our results are compatible with a late migration of the outermost planet into the resonant chain, when the giant planet pair already is in resonance.} | The interaction of planets with the protoplanetary disk of gas and dust in which they are born, allows them to exchange angular momentum and energy with it. This mechanism provides planets with the ability to change their orbital elements and migrate radially through the disk. Ever since this was described in the seminal work by \citet{GoldreichTremaine1980}, plenty of numerical and analytical studies have focused on determining migration rates in different regimes depending on the physical properties of planet and disk \citep[for a review, see][]{KleyNelson2012}. Even for the case of single planets, this is still an area of active research. Since such studies have revealed a wide range of possible migration rates, planets in a multi-planet system can migrate in a convergent fashion, for example if a planet migrates inward at a higher rate than an interior companion. The mutual gravitational interaction between the planets allows them to become captured into mean-motion resonances (MMR) that can have a major impact on the resulting orbital architecture. While most of the exoplanet systems that are known today are found to reside outside of exact resonances \citep{WinnFabrycky2015}, several systems are in a resonant state. Some cases of very compact resonant chains exist, as for example in the famous system around TRAPPIST-1 \citep{Luger2017}, where seven low-mass planets form such a configuration. While generally, convergent disk-driven migration has been shown to be an effective channel for the formation of orbital resonances, the modelling often uses damping prescriptions that are obtained by studying the migration in hydrodynamical models of single planets. However, for planet pairs that are made of planets in different mass regimes, for example an interior giant and an exterior super-Earth, it has been shown that disk-mediated perturbations can change the migration behaviour \citep[e.g.][]{PG2012,Baruteau2013}. The nearby red M-dwarf star Gliese 876 (GJ 876) hosts four known planets, that constitute one of the most prominent and dynamically interesting systems that we know of to date. A thorough review of the observational history of this planetary system, which spans two decades, is given in the introduction of \citet{NelsonB2016} and we just summarize some important characteristics of this system. The growing amount of radial-velocity (RV) observations from several instruments have subsequently revealed the first confirmed extrasolar mean-motion resonance and, with the discovery of the fourth planet, the only known instance of an extrasolar three-body Laplace resonance \citep{Rivera2010}, which is comprised of the three outer planets. Using dynamical orbital fits, the state of the system has been constrained with high precision to be engaged in a deep 4:2:1 MMR, exhibiting libration of the critical resonant angle $\varphi_\mathrm{L} = \lambda_\mathrm{c} - 3 \lambda_\mathrm{b} + 2\lambda_\mathrm{e}$ around zero with a low amplitude of about $30^\circ$\citep{NelsonB2016,Trifonov2017,Millholland2018}. However, it is still unclear whether the system is in a low-energy state of pure-double apsidal corotation (ACR), where the resonant angles always librate or a high-energy state of quasi ACR, where they switch between mostly libration and brief phases of circulation, since both are compatible with the most recent observations \citep{Millholland2018}. For reference, we show the parameters of the latest orbital fit by \citet{Millholland2018} in Table \ref{table:mill18}. \begin{table} \caption{Masses and orbital parameters of all four planets adopted from the best-fit model of \citet{Millholland2018}. More information and uncertainties are given in their Table 4.} \label{table:mill18} % \centering % \begin{tabular}{c c c c c} % \hline\hline % Planet & $M[M_\oplus]$ & $a\,[\mathrm{au}]$ & $e$ & $P\,[\mathrm{d}]$\\ % \hline % d & $7.55$ & 0.022 & 0.057 & 1.938\\ % c & $265.6$ & 0.136 & 0.257 & 30.10\\ % b & $845.2$ & 0.219 & 0.033 & 61.11\\ e & $15.6$ & 0.350 & 0.03 & 123.8\\ \hline % \end{tabular} \end{table} What makes this triplet of resonant planets even more interesting from a theoretical perspective is the fact that the two inner planets c and b, which were first discovered, are both giants with mass ratios of about $q_\mathrm{c} \simeq 2 \cdot 10^{-3}$ and $q_\mathrm{b} \simeq 6.5 \cdot 10^{-3}$ with respect to the star, that are equivalent to several Jupiter masses around a solar mass star. The outermost planet has a mass similar to Neptune and $q_\mathrm{e} \simeq 1.3 \cdot 10^{-4}$. Planets of such high masses are expected to significantly shape the disk, for example by opening deep gaps, which warrants detailed modelling of the migration process. Different to the regular Laplace resonance of the Galilean satellites where $\varphi_\mathrm{L}$ librates around $180^\circ$, the resonance in the system around GJ 876 is chaotic on relatively short time-scales. This behaviour has been recently analysed using analytical models and N-body integrations by \citet{Batygin2015}, who suggested that a purely dissipative evolution by migration in a laminar disk would not be able to produce the chaotic configuration and stochastic forcing by turbulence might be needed. In a follow-up study, \citet{Marti2016} identified two different regions they refer to as an inner and outer resonant region, where the latter is chaotic on decadal time-scales. This particular system has been recently used by \citet{Puranam2018} to study the chaotic eccentricity excitation of innermost planet, the super-Earth (d), which has a possibly low but finite value, in order to restrict its tidal dissipation. While this planet is not part of the resonant chain, it still experiences strong chaotic excitations from its companion planets. The fact that this system and its dynamical state are well-characterized makes it an interesting subject for theoretical studies that explore evolutionary paths that could lead to its formation. Just like the observations, such theoretical models have evolved with time and become more complex. By prescribing damping rates $\tau_{e,a}$ to the eccentricity and semi-major axis to mimic disk-driven migration in their N-body simulations, \citet{Lee2002} investigated the capture of the two giant planets into a 2:1 resonance. From their results, they constrained the ratio of these damping time-scales such that the obtained equilibrium eccentricities were compatible with the observations. First two-dimensional hydrodynamical models were presented by \citet{Snellgrove2001} and \citet{Kley2005}, who confirmed the possibility of capture into the observed resonance of the giant pair using both isothermal and radiative disk models. They also reported that the most massive planet in the system can excite substantial eccentricity in the disk, which was further investigated for a non-migrating giant in \citet{KleyDirksen2006}. However, only the most massive planet b was included in the active domain of their simulation, which could not provide sufficient damping rates to suppress the eccentricity growth of planet c, which orbited inside a central cavity. Further investigating this issue, \citet{Crida2008} considered the possibility of an inner disk that could provide eccentricity damping by directly interacting with planet c. They were able to match the observed resonant state of the giant pair quite accurately. Inferring damping time-scales from their hydrodynamical models allowed them to reproduce their findings using damped N-body models. Adopting a locally isothermal equation of state in both two- and three-dimensional models, \citet{Andre2016} also considered migrating giant planets and investigated the formation and maintainability of several orbital resonances. After previous works were focused on the interaction of giant planets, \citet{Podlewska2009,PG2012} modelled the migration of a super-Earth that orbits exterior to a giant planet, which is reminiscent of the planet-pair b-e in GJ 876. They found that the convergent migration of said planets can be halted or even reversed, depending on model parameters. They attributed this to the fact that, by means of the spiral waves that the giant planet launches, it is able to transport angular momentum and energy outward, depositing some fraction of it into the disk via shocks close to the super-Earth, which can then experience a modified positive torque. Of particular interest to our study, but at odds with the observation of this particular system is the fact, that most of their models cannot reproduce the 2:1 MMR of the outer planet pair around GJ 876. Divergent evolution of planet pairs via this mechanism has also been investigated by \citet{Baruteau2013}, mainly with the goal of explaining the observed resonant offset in the \textit{Kepler} planets. While above mentioned studies have considered the formation of MMR in a pair of adjacent planets, in this work, we aim to construct a self-consistent disk migration model of all three planets that are involved in the Laplace resonance. This will elucidate the role that previously described processes played during the formation and put constraints on the nature of the disk. This paper is structured as follows. In Section \ref{sec:setup} we describe the physical setup of our model and the governing equations, followed by an outline of the numerical methods that were used in Section \ref{sec:numerics}. In Section \ref{sec:ref} we show the results for a reference model that successfully produces a Laplace resonance. Models with varied disk parameters and boundary treatment are presented in Section \ref{sec:var}. We discuss our results in Section \ref{sec:discussion}. We summarize our findings and give an outlook on future work in Section \ref{sec:sum}. | \label{sec:sum} In this section we summarize our findings and give an outlook on possible directions of future studies. Using locally isothermal two-dimensional hydrodynamical simulations, we studied the migration of three planets of masses corresponding to those discovered around GJ 876. By varying physical disk parameters as well as the numerical modelling approach, we were able to find cases where the three-body resonance is formed and cases where the outer planet's migration is stalled outside of resonance. We now summarize the characteristics of models that produced similar results. \begin{itemize} \item[-] \textit{Successful formation of the Laplace resonance:}\\ Depending on the disk's viscosity and scale-height, the outer planet can open a partial gap and over long time-scales weaken the eccentricity-related repulsion by starving the giants' gap. In these cases the planet is able to reduce the disk eccentricity, possibly through a combination of starving the eccentricity-exciting resonances of the giant planet and providing eccentricity damping itself. As a result, associated torque variations are reduced and the outer planet is able to slowly migrate along the gap edge and settle into a 2:1 MMR with the giant planet, completing a 4:2:1 Laplace MMR that is similar to the observed state of the system. A late migration of the outermost planet into resonance seems likely from our findings, since a low disk-mass model also resulted in the formation of the Laplace resonance and the effects that lead to divergent evolution are expected to be weaker then. Disk parameters like a lower aspect ratio and viscosity, that favour the formation of the resonance are also compatible with this scenario. In a model where the giant planets are started closer to the star, the outer planet is unable to approach them, which suggests that the resonance around GJ 876 was formed in a wider configuration and became more compact over a possibly long time-scale until the disk dispersed. In our simulations, the Laplace resonance is in a state where only the inner planet pair is in ACR, while the outer pair is not. It remains to be explored whether this can change as the resonant chain becomes more compact until it matches the observed planet locations. \item[-] \textit{Stalled migration in an eccentric disk with a cavity:}\\ We find in many models, that the interaction of an outer lower mass planet with the wake driven by an interior giant planet and an eccentric gap edge can slow down or even halt the outer planet's inward migration. We showed that torque acting on the outermost planet shows oscillations on characteristic periods of the system and attribute them to a shift of resonances due to eccentricity of the disk and planets, as well as variations due to the spiral wakes of the interior giants. In order to disentangle different contributions, future studies should consider simpler situations and investigate their dependence on physical parameters of the disk and planets. Whenever the torque variations due to these effects are large, settling into resonance is prevented, even if the outer planet is stalled close to a commensurability. \item[-] \textit{Boundary conditions:}\\ Conducting many numerical experiments, we found that the choice of the inner BC and damping prescriptions used in its vicinity can have strong effects on the dynamical evolution of the system. Due to the high planetary masses and compactness of the system, we find a damping of the velocities close to the inner boundary to be necessary to avoid the formation of an artificial cavity, where the density is strongly reduced. A viscous evolution of the inner disk can still be accounted for by using velocity damping towards viscous outflow in combination with an outflow-only BC. \end{itemize} In our simplified approach, we modelled a two-dimensional, locally isothermal disk of constant aspect ratio. More realistically, the disk will be affected by irradiation and heating due to viscous dissipation and spiral shocks driven by giant planets, since not only angular momentum, but also energy is transferred \citep{Rafikov2016}. As has been shown for single planets already, radiative effects can substantially affect their migration rates, mainly by modifying the corotation torques \citep{Paardekooper2011}. A more realistic treatment of the disk thermodynamics is of special interest for the migration of an exterior lower-mass planet, since it might change the disk's eccentricity structure \citep{Teyssandier2016} and the interaction and shape of the giant planets' wake. In order to study the role of the disk and planet eccentricity in more detail, simpler situations should be considered, as for example a single planet migrating in an eccentric disk with a cavity. In order to find a simple description that could for example be used in N-body simulations, the dependence on the physical parameters of the disk and planet should be investigated in more detail in future work. In a three-dimensional model, additional effects could become important. Not only can the the growth-rates of eccentric modes be modified \citep{Teyssandier2016}, but the disk would be allowed to become warped and second-order inclination resonances would be available to the planets, which have been considered previously \citep{LeeThommes2009}. While we considered the planet masses to be fixed at their current values, mass growth by gas accretion adds additional pathways to the formation of this system. For example, a scenario where the planets grow already trapped in resonance could be imagined. However, given the potential starving of the inner planets that we observed, this seems rather unlikely. Due to the high masses of the giants, they must have accreted substantial amounts of gas during the gas-rich phase of the disk. The migration of a slightly eccentric planet in an eccentric disk warrants further investigation, since it is also of great interest for the migration of circumbinary planets. Since our study has revealed that the migration of the outer planet can be very different from a smooth, purely dissipative evolution, the observed torque variations might act in a similar way to turbulence in preventing settling into a regular state, which was proposed by \citet{Batygin2015}. | 18 | 8 | 1808.04223 |
1808 | 1808.06967_arXiv.txt | As first suggested by U. Fano in the 1940s, the statistical fluctuation of the number of pairs produced in an ionizing interaction is known to be sub-Poissonian. The dispersion is reduced by the so-called ``Fano factor," which empirically encapsulates the correlations in the process of ionization. In modeling the energy response of an ionization measurement device, the effect of the Fano factor is commonly folded into the overall energy resolution. While such an approximate treatment is appropriate when a significant number of ionization pairs are expected to be produced, the Fano factor needs to be accounted for directly at the level of pair creation when only a few are expected. To do so, one needs a discrete probability distribution of the number of pairs created $N$ with independent control of both the expectation $\mu$ and Fano factor $F$. Although no distribution $P(N|\mu,F)$ with this convenient form exists, we propose the use of the COM-Poisson distribution together with strategies for utilizing it to effectively fulfill this need. We then use this distribution to assess the impact that the Fano factor may have on the sensitivity of low-mass WIMP search experiments. | Following an ionizing particle interaction of deposited energy $E$, the mean number of pairs created is given by \begin{equation} \label{Np eq} \mu = \frac{E}{W(E)}, \end{equation} \noindent where $W$ is generally both energy and particle-type dependent and represents the mean energy needed to create either one electron-ion pair in liquid and gaseous detectors or one electron-hole pair in semiconductor devices \cite{sauli}. The actual number of pairs created $N$ is subject to statistical fluctuations, which limits the achievable energy resolution of any ionization measuring device to monoenergetic radiation. As first anticipated by Fano, the variance of these fluctuations $\sigma_N^2$ is lower than expected for a Poisson process by a factor $F$, known hereafter as the ``Fano factor," which is defined as \cite{uno2} \begin{equation} \label{Fdef eq} F = \frac{\sigma_N^2}{\mu}. \end{equation} By definition $F=1$ for a Poissonian process, whereas $F<1$ for ionization fluctuations. Experimentally measuring the Fano factor is challenging, as one needs to both strongly suppress and precisely quantify all sources of resolution degradation that do not arise from ionization fluctuations. In spite of this, measurements of $F$ have been carried out for a wide variety of materials including argon ($F \approx 0.23$), xenon ($F \approx 0.17$), silicon ($F \approx 0.16$), germanium ($F \approx 0.12$), and others \cite{scinprop,policarpo,owens,germanium}. Although these measurements set a fundamental upper limit on the resolution possible with these detector media, they do not provide more information about the actual probability distribution of the number of pairs created than its dispersion. While the latter can be predicted with Monte Carlo simulations of the processes involved in energy loss at a microscopic level \cite{gross1} (see Supplemental Material \cite{supp}), this approach is too computationally expensive to be practical for most applications. This is true for any scenario in which one needs to simulate the measurement of a signal that is not monoenergetic. In this case, one might think to fold the effect of the Fano factor into the overall energy resolution. While such an approach is appropriate at high energies, it is not valid when $\mu$ is small, in which case a more accurate treatment is necessary. To account for the Fano factor at the level of pair creation, one would require a discrete probability distribution $P \left( N | \mu,F \right)$ of the number $N$ of pairs created for any value of $\mu$ and $F$. Although there is no distribution with this exact convenient form, we propose the COM-Poisson distribution as a viable solution. It is a discrete distribution that allows for independent control of the mean and variance with two parameters, $\lambda$ and $\nu$. While these variables do not correspond to $\mu$ and $F$, we have developed a methodology to effectively translate $P \left( N | \lambda,\nu \right)$ into $P \left( N | \mu,F \right)$. \indent We believe the COM-Poisson distribution may provide a much needed tool in the area of low-mass dark matter research. Since new, popular models favor particle masses on the order of a few $\mathrm{GeV/c^2}$ or less \cite{Essig2013,ZUREK201491}, a growing cohort of direct detection experiments are now confronted with the issue of modeling ionization statistics at the single pair regime. This includes gaseous dark matter search experiments like NEWS-G \cite{Arnaud2018}, liquid noble experiments like DarkSide \cite{darkside1}, and solid-state experiments such as SuperCDMS, Edelweiss, DAMIC, and Sensei \cite{single,edelweiss,damic,sensei}. While these detector technologies differ in many ways, the requirements for modeling ionization statistics are essentially the same for each, and are fulfilled by the COM-Poisson distribution. What follows is a more detailed discussion about the problem of modeling ionization statistics (Sec.\ \ref{model sec}), the COM-Poisson distribution (Sec.\ \ref{COM sec}), and strategies for using it (Sec.\ \ref{using sec}). Finally, we use the COM-Poisson distribution to assess the potential impact of the Fano factor on the sensitivity of dark matter detection experiments in Sec.\ \ref{App sec}. | \label{discussion sec} In this work, we have proposed a novel approach and developed the strategies and tools to account for the Fano factor at the level of pair creation with the COM-Poisson distribution. By using it to assess the impact that the Fano factor may have on the sensitivity of low-mass WIMP search experiments, we have demonstrated both the need for modeling ionization statistics and the usefulness of the COM-Poisson distribution to do so. Although there is no physical motivation for the choice of COM-Poisson other than its apparent suitability for this application, we would like to stress that there is no comparable alternative to the proposed approach at the time of this publication. Additional rationale for using the COM-Poisson distribution can be found in \cite{supp}. To encourage and facilitate the use of this tool by others, we have provided free access to the look-up tables discussed in Sec.\ \ref{optimization sec} and the code to use them at \cite{website}. We have also provided code to produce the detection efficiency curves shown in Fig.\ \ref{Survive fig} as an illustrative example. The authors are open to providing assistance in using the tools developed and discussed in this work. | 18 | 8 | 1808.06967 |
1808 | 1808.01295_arXiv.txt | We compute Schwinger pair production rates at finite temperature, in the presence of homogeneous, concurrent electric and magnetic fields. Expressions are obtained using the semiclassical worldline instanton formalism, to leading order, for spin-$0$ and spin-$\frac{1}{2}$ particles. The derived results are valid for weak coupling and fields. We thereby extend previous seminal results in the literature, to coexistent electric and magnetic fields, and fermions. | The non-perturbative pair production of electrically and magnetically charged particles in the background of large field strengths has garnered much interest and study over the years. Sauter\Cit{Sauter:1931zz}, as well as Heisenberg and Euler\Cit{Heisenberg1936}, had speculated that sufficiently large electric fields could lead to spontaneous pair production of $e^+$- $e^-$. The notion was further sharpened and investigated comprehensively by Schwinger\Cit{Schwinger:1951nm}; deriving the imaginary part of the QED one-loop effective action. These results were then further generalised by various authors to diverse cases -- for instance, to extended objects such as magnetic monopoles\Cit{Affleck:1981ag}, spatial or temporal inhomogeneous fields\Cit{DunneWorldline, DunneFluctuationPre} and arbitrary gauge couplings\Cit{AFFLECK,Gould}, to cite a few examples (see for instance\Cit{Dunne:2004nc, Ruffini:2009hg} and references therein for a more complete discussion). Exact analytic expressions are known nevertheless only for a few special cases and extending investigations into hitherto unexplored regimes is an ongoing endeavour. The worldline path integral formalism has proven to be a potent method for perturbative and non-perturbative quantum field theoretic computations. The origins of the method may be traced to ideas by Fock\Cit{Fock:1937dy}, Nambu\Cit{Nambu:1950rs} and the Feynman worldline representation of one-loop effective actions\Cit{Feynman:1950ir,Feynman:1951gn}. The formalism was, for instance, leveraged to compute pair-production rates for magnetic monopoles at strong coupling\Cit{Affleck:1981ag,AFFLECK}. With the development of string theoretic techniques towards understanding gauge theory scattering amplitudes\Cit{Halpern:1977he,Halpern:1977ru,Polyakov:1987ez,Bern:1991aq,Strassler:1992zr}, the method found further resurgence and applications (see for example\Cit{Schubert:2001he} and related references); particularly, in our context, conveniently accommodating computations with large external fields\Cit{Schmidt:1993rk,Shaisultanov:1995tm,Adler:1996cja, Reuter:1996zm,Bastianelli:2004zp}. Among the pertinent extensions to non-perturbative pair production rates at zero temperature, are the inclusion of finite temperature corrections. This has received much attention in the literature\Cit{Dittrich, Cox:1984vf, Selivanov:1986tu, LoeweRojas,Elmfors:1993wj, Ganguly,Hallin,ELMFORS1995141,Gies:1998vt,Gies:1999vb,Gavrilov:2007hq, Gavrilov:2007ij,Kim:2010qq,King:2012kd, Brown:2015kgj, Medina,Gould,Gould:2018ovk}. There has been some discussion and disagreement in the literature though, over these thermal corrections, particularly in the constant electric field case lately\Cit{Brown:2015kgj,Medina,Gould,Gould:2018ovk}. Thermal corrections for this case was recently computed\Cit{Medina} and extended to arbitrary coupling\Cit{Gould,Gould:2018ovk}, using worldline path integral techniques. Our aim in this work is to extend these results to the case when there are homogeneous (spatially and temporally) electric and magnetic fields simultaneously present. We compute leading order thermal corrections, using worldline path integral techniques, to the non-perturbative vacuum decay rates when there are coexistent electric and magnetic fields. We work in a regime where the coupling constant is small, and the external fields are also relatively weak. As far as we know, these expressions have not been computed before in the literature. We will largely follow techniques developed in\Cit{Affleck:1981ag,DunneWorldline, DunneFluctuationPre,Schubert:2001he, Medina}. In the limit of vanishing temperature ($T\rightarrow 0$), one recovers the well-known results in literature\Cit{Schwinger:1951nm, Nikishov:1969tt, 1970SPhD...14..678B,Popov:1971iga, Daugherty:1976mg, Cho:2000ei, KimDONPageEparallelB}. When the magnetic field vanishes ($B\rightarrow 0$), in the case of scalar quantum electrodynamics (SQED), the results are seen to relapse into the known expressions for pure homogeneous electric fields, computed recently\Cit{Medina}. In quantum electrodynamics (QED), with fermions, we also obtain new expressions in the $B\rightarrow 0$ limit that complement these recent SQED results. It is well known that even at zero temperature ($T=0$), the presence of a magnetic field parallel to the electric field ($E \shortparallel B$), leads to interesting modifications to vacuum decay rates, relative to the pure electric field case. The vacuum decay rates, per unit volume, at $T=0$ for homogenous $E \shortparallel B$ are given by\Cit{Schwinger:1951nm, Nikishov:1969tt, 1970SPhD...14..678B,Popov:1971iga, Daugherty:1976mg, Cho:2000ei, KimDONPageEparallelB} \bea \label{eq:ebsppt0} \Gamma_{T=0, \text{\tiny{scalar}}}^{E \shortparallel B} &=& \sum_{k=1}^{\infty}\frac{(-1)^{k+1}q^{2}EB}{8\pi^{2}k\sinh(k\pi B/E)} \exp\Big{[}-\frac{m^{2}k\pi}{q E}\Big{]} \; , ~~~~\\ \Gamma_{T=0, \text{\tiny{fermion}}}^{ E \shortparallel B} &=& \sum_{k=1}^{\infty}\frac{q^{2}EB\coth(k\pi B/E)}{4\pi^{2}k} \exp\Big{[}-\frac{m^{2}k\pi}{q E}\Big{]} \nn \; . \eea Here, $m$ and $q$ are the mass and electric charge of the particle under consideration. Note that in addition to the usual enhancement due to extra degrees of freedom in the spin-$\frac{1}{2}$ case, the vacuum decay rates in the fermion case may be further enhanced, relative to the scalar case, when $B > E$. Note also that, for any homogeneous $\vec{E}^{'}$ and $\vec{B}^{'}$ fields, for which the Lorentz invariant $\vec{E}^{'}\cdot \vec{B}^{'} \neq 0$, one may go to a frame of reference with boost ($\vec{\upsilon}$) given by\Cit{Landau:1982dva} \be \frac{\vec{\upsilon}}{1+\lvert\vec{\upsilon}\rvert^2}=\frac{\vec{E}^{'}\times \vec{B}^{'}}{\lvert \vec{E}^{'}\rvert^2+\lvert \vec{B}^{'}\rvert^2}\; , \ee where the transformed fields ($\vec{E}$ and $\vec{B}$) are parallel to each other. This is the so-called centre-of-field frame. Since the vacuum decay rate per unit volume is a Lorentz invariant, one may conveniently compute it in this centre-of-field frame. The formulas for homogeneous $E \shortparallel B$ are therefore potentially of wide applicability. For homogeneous fields with $\vec{E}^{'}\cdot \vec{B}^{'} = 0$, but the fields not equal in magnitude, a reference frame may be found where the transformed field is purely electric or magnetic\,\cite{Landau:1982dva}. In this latter scenario, the relevant expressions are those of single field Schwinger pair production. Apart from being of significant theoretical interest, scenarios with parallel electric and magnetic fields are also relevant in various astrophysical systems. For instance, it is believed that neutron stars such as pulsars have strong electrical fields parallel to the magnetic field in their polar vacuum gap regions\Cit{2004hpa..book.....L}. Neutron star surface temperatures are expected to reach $\sim 10^5\,\mathrm{K}$. Non perturbative production of exotic states such as millicharged particles, which may form a component of dark matter, may occur in these vacuum gap regions and provide hitherto unknown constraints on these states\Cit{Korwar:2017dio} (in the context of constraints from non-perturbative production, in pure $E$ or $B$ fields, also see \Cit{Gies:2006hv} for millicharged particle bounds from accelerator cavities, and \Cit{Hook:2017vyc, Gould:2017zwi} for bounds on magnetic monopoles). These settings also, therefore, make the results phenomenologically very relevant. In \Sec{sec:sqed}, we discuss the derivation in the case of SQED. Towards the exposition of necessary techniques and to fix notations, we re-derive the known zero temperature result for the case of $E \shortparallel B$ using worldline instantons, before presenting the main result for finite temperature. Then, in \Sec{sec:qed}, we consider QED. Results are presented for spin-$\frac{1}{2}$ particles in the zero temperature and finite temperature cases. The finite temperature SQED and QED Schwinger pair production results, for $E \shortparallel B$, are new and readily generalise earlier seminal results in the literature\Cit{Brown:2015kgj, Medina,Gould}. Even for the zero temperature cases, to the best of our knowledge, this is the first time that an explicit and complete derivation is being presented for vacuum decay rates, when $E \shortparallel B$, using worldline instanton techniques. We summarise our main results, shortcomings of the derivations and future directions in \Sec{sec:conc}. | {\label{sec:conc}} The worldline path integral formalism provides a powerful and systematic way to compute nonperturbative vacuum decay rates in various situations. In this work, we computed leading order thermal corrections to vacuum decay rates, in SQED and QED, for the case of homogeneous, coexistent electric and magnetic fields. Apart from its theoretical importance, the results are relevant in astrophysical settings where large electric and magnetic fields may coexist in a thermal environment. There are a few natural avenues to follow up on that were outside the scope of the present study. The Gaussian approximation to the fluctuation prefactor is inadequate, leading to spurious singularities at thermal thresholds, and one should include higher order terms to potentially mitigate this. This is challenging even in the zero temperature case, but based on the hard thermal loop framework\Cit{Braaten:1991gm,Braaten:1992gd}, it has been argued that such spurious singularities may be softened and the result correctly interpreted\Cit{Medina}. Explicit calculation of these higher order terms beyond the Gaussian approximation would shed more light on the analytic structure of the terms at these thresholds. Another subtle point to note is that, even at zero temperature, the vacuum decay rate is not technically the same as the average, particle pair production rate\Cit{Nikishov:1970br, Cohen:2008wz}. In the zero temperature case, it may be shown that the physical observable--the mean pair production rate--is just the first term in the series for the vacuum decay rate\Cit{Nikishov:1970br}. Hence, for weak fields, the distinction is mostly pedantic. The thermal vacuum decay rates we compute are therefore expected to closely match the actual particle pair production rates for weak fields, but a more careful calculation is required to make the correspondence clear and rigorous. It would also be appealing to have a better physical understanding of the various results and reach a consensus on the remaining disagreements in the literature\Cit{Brown:2015kgj,Medina,Gould,Gould:2018ovk}. Doing away with the assumption of relatively weak fields and extending the study to arbitrary coupling strengths would also be pertinent, as well as incorporating modifications due to field inhomogeneities. | 18 | 8 | 1808.01295 |
1808 | 1808.00770_arXiv.txt | In this article, we present a simulator conceived for the conceptual study of an AO-fed high-contrast coronagraphic imager. The simulator implements physical optics: a complex disturbance (the electric field) is Fresnel-propagated through any user-defined optical train, in an end-to-end fashion.\\The effect of atmospheric residual aberrations and their evolution with time can be reproduced by introducing in input a temporal sequence of phase screens: synthetic images are then generated by co-adding instantaneous PSFs. This allows studying with high accuracy the impact of AO correction on image quality for different integration times and observing conditions.\\In addition, by conveniently detailing the optical model, the user can easily implement any coronagraphic set-up and introduce optical aberrations at any position. Furthermore, generating multiple images can allow exploring detection limits after a differential post-processing algorithm is applied (e.g. Angular Differential Imaging).\\The simulator has been developed in the framework of the design of SHARK-NIR, the second-generation high contrast imager selected for the Large Binocular Telescope. | \label{sec:intro} % SHARK-NIR is a coronagraphic camera. Since coronagraphs deal with diffraction of light, it is necessary to operate in the framework of wave-propagation physics: given the mathematical complexity of this theory, the most common approach is to make use of numerical simulations. The test bench is written in IDL language and it is developed to assess the coronagraphic performance of the camera.\\Section \ref{sec:SC} describes the simulator, with a brief insight into critical aspects such as implementation of AO correction and NCPA. Section \ref{sec:coro} introduces the coronagraphic techniques implemented in the test bench and identified as possible candidates for SHARK-NIR. Finally, section \ref{sec:results} shows some of the possible studies that can be performed with the simulator. | \label{sec:CON} In this paper we have presented a Fresnel simulator developed in the framework of the SHARK-NIR project, the second generation high-contrast imager of the Large Binocular Telescope. The simulator is conceived as a test bench: it implements several coronagraphic techniques and can generate images in a wide range of operative conditions. Several sources of optical aberrations are introduced, with particular care on realistically reproducing the LBT environment. The examples of possible applications shown here are part of a more comprehensive tradeoff study that led to the definition of an optimal suite of coronagraphic designs for the instrument. \newpage | 18 | 8 | 1808.00770 |
1808 | 1808.05363_arXiv.txt | By starting with a two-fields model in which the fields and their derivatives are nonminimally coupled to gravity, and then by using a conformal gauge, we obtain a model in which the derivatives of the canonically normalized field are nonminimally coupled to gravity. By adopting some appropriate functions, we study two cases with constant and E-model nonminimal derivative coupling, while the potential in both cases is chosen to be E-model one. We show that in contrary to the single field $\alpha$-attractor model that there is an attractor \textit{point} in the large $N$ and small $\alpha$ limits, in our setup and for both mentioned cases there is an attractor \emph{line} in these limits that the $r-n_{s}$ trajectories tend to. By studying the linear and nonlinear perturbations in this setup and comparing the numerical results with Planck2015 observational data, we obtain some constraints on the free parameter $\alpha$. We show that by considering the E-model potential and coupling function, the model is observationally viable for all values of $M$ (mass scale of the model). We use the observational constraints on the tensor-to-scalar ratio and the consistency relation to obtain some constraints on the sound speed of the perturbations in this model. As a result, we show that in a nonminimal derivative $\alpha$-attractor model, it is possible to have small sound speed and therefore large non-Gaussianity.\\ {\bf PACS}: 98.80.Cq , 98.80.Es\\ {\bf Key Words}: Cosmological Perturbations, Non-Gaussianity, Nonminimal Derivative Coupling, $\alpha$-Attractor, Observational Constraints. | Considering a single canonical scalar field (inflaton) responsible for cosmological inflation in early universe, is a simplest way to solve some problems of the standard model of cosmology. To have enough e-folds number or equivalently enough exponential expansion of the universe, the inflaton should rolls slowly down toward the minimum of its potential. In this simple inflation model, an adiabatic, scale invariant and gaussian mode of the primordial perturbations is dominant~\cite{Gut81,Lin82,Alb82,Lin90,Lid00a,Lid97,Rio02,Lyt09,Mal03}. However, it seems that in the future, with the advancement of technology, we should be able to detect the non-Gaussian distributed modes of the perturbations. Also, some extended models of inflation predict scale dependent and non-Gaussian features of the primordial perturbations~\cite{Mal03,Bar04,Che10,Fel11a,Fel11b,Noz12,Noz13a,Noz13b,Noz13c}. In this regard, in studying the inflation, the models predicting the non-Gaussian perturbations are really interesting. It is possible the scalar field responsible for cosmological inflation to be the Higgs boson. It is proposed that to adopt the Higgs boson as the inflaton, we should consider a nonminimal coupling between its derivatives and Einstein tensor~\cite{Ame93,Ger10}. In this case, the friction is enhanced gravitationally at higher energies because of the presence of nonminimal derivative coupling. This means that the friction of an inflaton rolling down its own potential increases significantly, allowing occurrence of the slow-roll inflation even with steep potentials. The models with nonminimal derivatives coupling are capable to solve the unitary violation problem during inflation. In these models, unitarity is not violated up to the quantum gravity scales and also, quantum gravity regime is avoided during Inflation~\cite{Ger10,Ger12a}. Note that, in~\cite{Ger10} it has been shown that to trust the effective inflationary description, the curvature should be much smaller than the Planck scale. Considering the relation $R=6(\dot{H}+2H^2)$ and the fact that during inflation era $H$ is nearly constant and $\dot{H}$ is very small, we have $R\simeq 12H^2$. So we can say that, to avoid the unitarity problem during inflation, $H$ should be much smaller than Planck mass and this is the unitarity bound. In the nonminimal derivative model this bound is not violated~\cite{Ger10,Ger12a,Ger12b}. Also, in these models the perturbations are somewhat scale dependent and it is possible to have non-Gaussian distributed perturbations. We refer to~\cite{Tsu12,Sar10,Noz16a,Noz16b} for some works on the issue of nonminimal derivatives in the early time accelerating expansion of the universe as well as the late time cosmic dynamics. Recently, the idea of ``cosmological attractor'' in the models describing the cosmological inflation has attracted a lot of attention. The conformal attractors~\cite{Kal13a,Kal13b} and $\alpha$-attractors~\cite{Kai14,Fer13,Kal13c,Kal14} models are some models which incorporates the idea of the cosmological attractors. In Refs.~\cite{Cec14,Kal13d,Kal14b,Lin15,Jos15a,Jos15b,Kal16,Sha16,Odi16,Ras18} some more details on the issue of $\alpha$-attractors have been studied. The conformal attractor model has the universal prediction in the large e-folds number ($N$) for the scalar spectral index and tensor-to-scalar ratio as $n_{s}=1-\frac{2}{N}$ and $r=\frac{12}{N^{2}}$, respectively. The $\alpha$-attractor models are divided into two categories named E-model and T-model, according to the adopted potentials. The E-model corresponds to the following potential \begin{equation} \label{eq1} V=V_{0}\Big[1-\exp\big(-\sqrt{\frac{2\kappa^{2}}{3\alpha}}\phi\big)\Big]^{2n}\,, \end{equation} and the T-model is characterized by a potential as \begin{equation} \label{eq2}V=V_{0}\tanh^{2n}\Big(\frac{\kappa\phi}{\sqrt{6\alpha}}\Big)\,. \end{equation} In these potentials, $V_{0}$, $n$ and $\alpha$ are some free parameters. The prediction of the scalar spectral index in the $\alpha$-attractor model is similar to the prediction in the conformal attractor ones as $n_{s}=1-\frac{2}{N}$ in small $\alpha$ and large $N$ limits. However, it predicts the tensor-to-scalar ratio as $r=\frac{12\alpha}{N^{2}}$ in small $\alpha$ and large $N$ limits, which is somewhat different from the one predicted by the conformal attractor model. In this paper, we are going to study a nonminimal derivatives model in which both the potential and nonminimal derivatives coupling function are E-model type. Actually, the author of Ref.~\cite{Tsu12} has studied the nonminimal derivatives model in which the coupling function is a constant. He has adopted several types of potential such as $\phi^{2}$, $\phi^4$, exponential and so on. Our attention here is on $\phi^{2}$ potential. In Ref.~\cite{Tsu12}, it has been shown also that the nonminimal derivatives model with $\phi^{2}$ potential is consistent with joint data of WMAP7~\cite{Kom11}, BAO~\cite{Per10}, and HST~\cite{Rie09} for $N=50$, $60$ and $70$. However, if we compare the results with Planck2015 TT, TE, EE+lowP data~\cite{Ade15a} the model is observationally viable just for some values of $M$ (the mass scale of the nonminimal derivative coupling). We are going to check that by considering E-type potential and coupling function, whether the numerical results of the model are consistent with Planck2015 TT, TE, EE+lowP data background for all values of $M$ or not. In this regard, we follow Refs.~\cite{Kal13a,Kal13b} and consider a model with two real scalar fields $\psi$ and $\varphi$. The nonminimal action written in Ref.~\cite{Kal13a} is conformal invariant. The authors in this reference have used a $SO(1,1)$ conformal gauge (named rapidity gauge) as \begin{equation}\label{eq3} \psi^{2}-\varphi^{2}=6 \end{equation} which represents a hyperbola. By using a canonically normalized field as \begin{equation}\label{eq4} \psi=\sqrt{6}\cosh (\frac{\phi}{\sqrt{6}}) \,\,, \quad \varphi=\sqrt{6}\sinh (\frac{\phi}{\sqrt{6}}) \end{equation} they were able to eliminate the nonminimal terms in the action and transform the nonminimal action to the minimal one accordingly. Actually, by adopting several types of the potential terms, they have obtained the models corresponding to dS/AdS space, T-model of chaotic inflation and Starobinsky model of inflation~\cite{Sta80,Sta83,Whi84}. In our setup, both the fields and their derivatives are nonminimally coupled to gravity. To eliminate the nonminimal coupling ( not the ``nonminimal derivatives coupling'') we use the gauge (\ref{eq3}) and also rewrite equation (\ref{eq4}) as \begin{equation}\label{eq5} \psi=\sqrt{6}\cosh (\frac{\phi}{\sqrt{6\alpha}}) \,\,, \quad \varphi=\sqrt{6}\sinh (\frac{\phi}{\sqrt{6\alpha}}) \end{equation} where the free parameter $\alpha$ has been included which leads us to E-model $\alpha$-attractor. Actually, $\alpha$ is inversely proportional to the curvature of the inflaton K\"{a}hler manifold~\cite{Kal13c}. By using this field, we re-parameterize the two fields model and convert it to a one-filed model with nonminimal derivatives coupling shown by the function ${\cal{F}}$. We study two cases as ${\cal{F}}=const.$ and ${\cal{F}}={\cal{F}}(\phi)$ and then we study cosmological inflation and perturbations in this setup. We show that in both cases there is an attractor \textit{line} in large $N$ and small $\alpha$ limits which the $r-n_{s}$ trajectories tend to. Note that, as we said, in the the single field $\alpha$-attractor model there is an attractor \textit{point} in these limits. In Ref.~\cite{Noz17} it has been shown that in the Gauss-Bonnet $\alpha$-attractor model also, there is an attractor \textit{point} in the mentioned limits. Indeed, the presence of line instead of point is because of considering the nonminimal derivatives coupling which causes the scalar spectral index of this model to be a functions of $\alpha$ and $M$. For $M\rightarrow \infty$, we recover the attractor point in usual $\alpha$-attractor models. In section 2, we study the inflation in this NMDC $\alpha$-attractor model and obtain the background equations of the model. In section 3, by using the ADM formalism, we study the linear perturbations in this model. In this section, we obtain some expressions for the scalar spectral index and tensor-to-scalar ratio and compare the results with Planck2015 observational data to test the observational viability of the model. In this regard, we obtain some constraints on parameter $\alpha$. In section 4, we study the non-linear perturbations and non-Gaussian features of the primordial perturbations. By using the relations between the tensor-to-scalar ratio and sound speed, and using the allowed values of the tensor-to-scalar ratio (from $95\%$ CL of Planck2015 TT, TE, EE+lowP data), we obtain constraints on the sound speed in this model. These constraints show that it is possible to have large non-Gaussianity in this model. In section 5 we present a summary of our work. | The aim of this paper was to study the non-minimal derivative model in the context of $\alpha$-attractor. In this regard, we have considered a two-fields action in which both the fields and their derivatives are nonminimally coupled to the gravity. We have used the conformal gauge to convert the two-fields model into a one field which its derivative (and not the field itself) is nonminimally coupled to the gravity. After obtaining the the background equations of this NMDC model, we have studied the linear perturbations theory. By adopting the ADM formalism, we have studied both the scalar and tensor perturbations in this setup and found some important perturbation parameters like as the power spectrum, scalar spectral index and tensor-to-scalar ratio. After that, by introducing the types of the general functions $\hat{{\cal{F}}}(\psi,\varphi)$ and $U(\psi,\varphi)$, and using the conformal gauge, the NMDC function and potential have been obtained. In this paper, two types of NMDC function have been considered; ${\cal{F}}=const.$ and ${\cal{F}}={\cal{F}}(\phi)$. We have shown that in both cases there is an attractor \textit{line} in large $N$ and small $\alpha$ limits which the $r-n_{s}$ trajectories tend to, while in the the single field $\alpha$-attractor model there is an attractor \textit{point} in these limits. In fact, the presence of attractor line instead of attractor point is due to the nonminimal derivatives coupling that causes the scalar spectral index of this model to be a functions of two parameters, $\alpha$ and $M$. In the limit $M\rightarrow \infty$, one recovers the attractor \emph{point} as is usual in $\alpha$-attractor models. We have numerically studied the linear perturbation in two cases with constant NMDC coupling and E-model NMDC coupling (in both cases the potential is considered to be E-model). With these choices, the power spectrum of the perturbations in NMDC model can get the observationally viable value (almost $2.4\times10^{-9}$) in some ranges of $\chi$ and $\alpha$. Also, in this model $r-n_{s}$ plane lies in the background of Planck2015 TT, TE, EE+lowP data for all values of $M$ (or $\chi$) and in some ranges of $\alpha$, for both $N=60$ and $N=70$. We have obtained some constraints on the free parameter $\alpha$ which lead to the values of the scalar spectral index and tensor-to-scalar ratio which are consistent with $68\%$ and $95\%$ CL of the Planck2015 TT, TE, EE+lowP data for $r-n_{s}$ distribution. By studying the nonlinear perturbation and three point correlation functions, we have studied the non-Gaussian feature of the primordial perturbations. In this regard we have obtained the amplitude of the equilateral configuration of the non-Gaussianity in terms of the sound speed. The sound speed is related to the tensor-to-scalar ratio via the consistency relation. The tensor-to-scalar ratio is constrained by using the $95\%$ CL of the Planck2015 TT, TE, EE+lowP data for $r-n_{s}$ distribution. By using the constraints on the tensor-to-scalar ratio, we have obtained some constraints on the values of the sound speed in this model. In this regard, we have shown that it is possible to have small sound speed leading to the large non-Gaussianity in this NMDC $\alpha$-attractor model.\\ {\bf Acknowledgement}\\ We would like to appreciate the referee for very insightful comments that have boosted the quality of the paper considerably. This work has been supported financially by Research Institute for Astronomy and Astrophysics of Maragha (RIAAM) under research project number 1/5237-97. | 18 | 8 | 1808.05363 |
1808 | 1808.05964_arXiv.txt | The near-infrared background between 0.5 $\mu$m to 2 $\mu$m contains a wealth of information related to radiative processes in the universe. Infrared background anisotropies encode the redshift-weighted total emission over cosmic history, including any spatially diffuse and extended contributions. The anisotropy power spectrum is dominated by undetected galaxies at small angular scales and diffuse background of Galactic emission at large angular scales. In addition to these known sources, the infrared background also arises from intra-halo light (IHL) at $z < 3$ associated with tidally-stripped stars during galaxy mergers. Moreover, it contains information on the very first galaxies from the epoch of reionization (EoR). The EoR signal has a spectral energy distribution (SED) that goes to zero near optical wavelengths due to Lyman absorption, while other signals have spectra that vary smoothly with frequency. Due to differences in SEDs and spatial clustering, these components may be separated in a multi-wavelength-fluctuation experiment. To study the extent to which EoR fluctuations can be separated in the presence of IHL, extra-galactic and Galactic foregrounds, we develop a maximum likelihood technique that incorporates a full covariance matrix among all the frequencies at different angular scales. We apply this technique to simulated deep imaging data over a 2$\times$100 deg$^2$ sky area from 0.75 $\mu$m to 5 $\mu$m in 9 bands and find that such a ``frequency tomography'' can successfully reconstruct both the amplitude and spectral shape for representative EoR, IHL and the foreground signals. | The optical and infrared background traces nucleosynthesis in stars and radiation from black holes throughout cosmic history. In addition to sources in our Galaxy, the absolute infrared background intensity is composed of diffuse sources of emission~\citep{2002MNRAS.336.1082S}, such as intra-halo light (IHL)~\citep{2012arXiv1210.6031C, 2014Sci...346..732Z} and faint galaxies present during the epoch of reionization (EoR). Instead of absolute intensity, which can be easily contaminated by Galactic emission, the near-infrared background may be studied using spatial fluctuations or anisotropies~\citep{2016RSOS....350555C}. These have been measured in broad continuum bands by a number of groups~\citep{2011ApJ...742..124M, 2005Natur.438...45K, 2012arXiv1210.6031C, 2012ApJ...753...63K, 2015NatCo...6E7945M, 2014Sci...346..732Z}. The fluctuation amplitude, which is robustly consistent across these measurements, exceeds that expected from the large-scale clustering of known galaxy populations~\citep{2012ApJ...752..113H}. The EoR signal is associated with the first collapsed objects that formed and produced energetic ultraviolet (UV) photons that reionized the surrounding hydrogen gas. In theoretical models~\citep{2012ApJ...756...92C, 2012ApJ...750...20F}, the expected amplitude of this component is several orders of magnitude below the level from faint galaxies in the more nearby universe. The emission of the ionizing photons during reionization is expected to peak between 0.9 and 1.1 $\mu$m today, if the implied optical depth is consistent with Planck result~\citep{2016A&A...596A.108P} and the reionization occurred around $z\sim$ 7 to 9. Also, it is damped quickly shortward of $\sim0.8\, \mu \rm{m}$~\citep{2004ApJ...606..683S, 2003MNRAS.339..973S} due to Lyman absorption. While the amplitude is small, component separation is possible through the unique spectral dependence afforded by the Lyman drop-out signature, similar to the Lyman drop-out signature used to identify bright galaxies present at $z > 6$ during reionization. Therefore, spatial fluctuations of the infrared (IR) background centered around 1 $\mu$m provide a way to discriminate the signal generated by galaxies present during reionization from those at lower redshifts, based on the strength of the drop-out signature in the fluctuations measured in different bands. In a recent study~\citep{2015NatCo...6E7945M}, a model with different IR fluctuations was constrained by multi-wavelength power spectra in five bands conducted with the Hubble Space Telescope (HST)/Advanced Camera for Surveys (ACS) and Wide Field Camera 3 (WFC3) data from the Cosmic Assembly Near-IR Deep Extragalactic Legacy Survey (CANDELS). Instead of only using auto-correlations of the HST measurements that are largely limited by the number of broad bands~\citep{2015NatCo...6E7945M}, in this work, we develop a novel component separation approach with full covariance information, including cross correlations of fluctuations across multiple frequency bands, to extract both spectral and spatial information for various astrophysical components. The cross correlations of spatial fluctuations between different wavelengths can break degeneracies with model parameters of other components, thereby providing crucial information about the EoR. This type of component separation will be crucial for upcoming measurements that will be performed with more than five spectral channels, such as the Cosmic Infrared Background ExpeRiment-2 (CIBER-2)\footnotemark[1]\footnotetext[1]{https://cosmology.caltech.edu/projects/ciber}~\citep{2014Sci...346..732Z,2016SPIE.9904E..4JS} and the proposed all-sky Spectro-Photometer for the History of the Universe, Epoch of Reionization, and Ices Explorer (SPHEREx) \footnotemark[2]\footnotetext[2]{http://spherex.caltech.edu/}, both of which can provide ideal data sets to measure the IR background fluctuations and study the EoR~\citep{2016arXiv160607039D, 2014arXiv1412.4872D}. In this work, we focus on the SPHEREx experiment, which has one observing straitening that will produce two surveys of different depths: an all-sky shallow survey and two deep surveys near the ecliptic poles. SPHEREx provides measurements in a total of 96 bands from 0.75 to 5 $\mu$m, and we only focus on fluctuations measured by combining those narrow-band images to nine broad-band images -- 0.8, 0.9, 1.025, 1.2, 1.5, 2.0, 2.65, 3.5, 4.5 $\mu$\rm{m} -- over a 2$\times$100 deg$^2$ sky area from the deep survey. Even with the compressed bands, the component separation scheme starts to become computationally challenging as the total number of parameters grows very fast when all the cosmic infrared background (CIB) sources have to be modeled at $N$ broad bands. This paper is organized as follows: In Section 2 we summarize various CIB components and the methodology related to component separation is introduced in Section 3; in Section 4 we apply the component-separation method to the simulated SPHEREx data; we conclude with a summary in Section 5. Unless otherwise noted we assume a cosmological model consistent with latest Planck CMB measurements~\citep{2016A&A...594A..13P}.\\ | In this paper, we have calculated all the covariant power spectra among the $\ihl$, $\lz$ and $\eor$ components using the halo-model formalism. We create consistent DGL and shot noise models for SPHEREx and generate $\ihl$, $\lz$ and $\eor$ models from theoretical predictions. Mock CIB auto- and cross-power spectra at nine broad bands \mbox{--} 0.8, 0.9, 1.025, 1.2, 1.5, 2.0, 2.65, 3.5, 4.5 $\mu$\rm{m} are made, and the component separation method is applied to separate EoR signal from IHL and other signals with simulated data expected from the planned SPHEREx project. Using simulations, we find that the component separation procedure constructed by a maximum likelihood function with a full covariance among different wavelengths can successfully reconstruct any component in the CIB fluctuations without introducing any significant biases into each reconstructed component. From parameter samples of the Monte Carlo Markov chain, power-spectrum uncertainties are determined for all the auto-power spectra of $\ihl$, $\lz$ and EoR. The spectrum of EoR fluctuations measured by SPHEREx is constructed at a broad range of wavelengths, indicating that a $>5\sigma$ level detection significance can be reached for a wide range of input models and assumptions about the measurement. The component separation algorithm also passed a series of tests related to model assumptions and observational parameter changes such as the number and location of observing bands, verifying that it is generally robust to such variations. \newpage | 18 | 8 | 1808.05964 |
1808 | 1808.09467_arXiv.txt | We present kinematics of 35 highly \rpro-enhanced ([Eu/Fe] $\geq +$0.7) metal-poor ($-$3.8 $<$ [Fe/H] $< -$1.4) field stars. We calculate six-dimensional positions and velocities, evaluate energies and integrals of motion, and compute orbits for each of these stars using parallaxes and proper motions from the second \textit{Gaia} data release and published radial velocities. All of these stars have halo kinematics. Most stars (66\%) remain in the inner regions of the halo ($<$~13~kpc), and many (51\%) have orbits that pass within 2.6~kpc of the Galactic center. Several stars (20\%) have orbits that extend beyond 20~kpc, including one with an orbital apocenter larger than the Milky Way virial radius. We apply three clustering methods to search for structure in phase space, and we identify eight groups. No abundances are considered in the clustering process, but the [Fe/H] dispersions of the groups are smaller than would be expected by random chance. The orbital properties, clustering in phase space and metallicity, and lack of highly \rpro-enhanced stars on disk-like orbits indicate that such stars likely were accreted from disrupted satellites. Comparison with the galaxy luminosity-metallicity relation suggests $M_{V} \gtrsim -9$ for most of the progenitor satellites, characteristic of ultra-faint or low-luminosity classical dwarf spheroidal galaxies. Environments with low rates of star formation and Fe production, rather than the nature of the \rpro\ site, may be key to obtaining the [Eu/Fe] ratios found in highly \rpro-enhanced stars. | \label{intro} The heaviest elements found in many metal-poor stars were produced by the rapid neutron-capture process (\rpro) in earlier generations of stars. Work by \citet{gilroy88} demonstrated that genuine differences exist in the overall levels of enhancement of \rpro\ elements relative to Fe in metal-poor stars. The recognition of the highly \rpro-enhanced star \object[BPS CS 22892-052]{CS~22892-052} by \citet{sneden94} in the HK Survey of \citet{beers92} erased any lingering doubt about the inhomogeneous distribution of \rpro\ elements in the environments where metal-poor stars formed. Stars that exhibit Eu/Fe ratios at least 10 times higher than in the Sun, like \object[BPS CS 22892-052]{CS~22892-052}, comprise only a small fraction ($\approx$~3\%; \citealt{barklem05heres}) of all metal-poor field stars, and none are known to be physically associated with each other \citep{roederer09a}. The environmental impact of the \rpro---expressed through the occurrence frequency, distribution, and enhancement levels of \rpro\ elements in stars---can help associate \rpro\ abundance patterns with their nucleosynthetic origins. Observations of the kilonova associated with gravitational wave event GW170817 \citep{abbott17prl,abbott17multimessenger} provide the most direct confirmation that neutron-star mergers are a site capable of producing heavy elements by \rpro\ nucleosynthesis (e.g., \citealt{cowperthwaite17,drout17,kasen17,tanvir17}). The occurrence frequency and level of \rpro\ enhancement of stars in dwarf galaxies supports this conclusion (e.g., \citealt{ji16nat,safarzadeh17ret2,tsujimoto17}), although those results alone cannot exclude rare classes of supernovae as an additional site (e.g., \citealt{tsujimoto15mrsn,beniamini16b}). Chemical evolution models \citep{cote18rpro} and simulations \citep{naiman18} can help generalize this result to \rpro\ production in the Milky Way. The $^{244}$Pu abundance in deep-sea sediments, which can be used to infer the content of this \rpro-only isotope in the ISM, also points to rare \rpro\ events like neutron-star mergers \citep{hotokezaka15,wallner15}. We lack similar, direct knowledge of the birth environments of highly \rpro-enhanced stars in the Milky Way halo field. An increasing number of these stars are now known (e.g., \citealt{hansen18}, and other ongoing work by the \textit{R}-Process Alliance). Their proximity to the Sun permits detailed abundance inventories to be derived from optical, ultraviolet, and near-infrared spectra (e.g., \citealt{sneden98,roederer12d,afsar16}). Five-parameter astrometric solutions (parallax, right ascension, declination, proper motion in right ascension, proper motion in declination) are now available for many of these stars in the second data release of the \textit{Gaia} mission (DR2; \citealt{lindegren18}). Line-of-sight velocities based on high-resolution optical spectroscopy are also available for these stars. The full space motion of each star can be reconstructed from these six parameters once a Galactic potential is adopted. We use these data to examine the kinematic properties of a large sample of highly \rpro-enhanced field stars for the first time. We present our sample of highly \rpro-enhanced field stars in Section~\ref{sample}, and we present their astrometric and velocity data in Section~\ref{kinematicdata}. We describe our calculations of the kinematics in Section~\ref{orbits}. We discuss the implications of these calculations in Section~\ref{discussion}, and we summarize our conclusions in Section~\ref{conclusions}. Throughout this work, we adopt the standard nomenclature:\ for elements X and Y, [X/Y] is the abundance ratio relative to the Solar ratio, defined as $\log_{10} (N_{\rm X}/N_{\rm Y}) - \log_{10} (N_{\rm X}/N_{\rm Y})_{\odot}$. | 18 | 8 | 1808.09467 |
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1808 | 1808.02858_arXiv.txt | Finite size effects in a neutron star merger are manifested, at leading order, through the tidal deformabilities of the stars. If strong first-order phase transitions do not exist within neutron stars, both neutron stars are described by the same equation of state, and their tidal deformabilities are highly correlated through their masses even if the equation of state is unknown. If, however, a strong phase transition exists between the central densities of the two stars, so that the more massive star has a phase transition and the least massive star does not, this correlation will be weakened. In all cases, a minimum deformability for each neutron star mass is imposed by causality, and a less conservative limit is imposed by the unitary gas constraint, both of which we compute. In order to make the best use of gravitational wave data from mergers, it is important to include the correlations relating the deformabilities and the masses as well as lower limits to the deformabilities as a function of mass. Focusing on the case without strong phase transitions, and for mergers where the chirp mass ${\cal M}\le1.4M_\odot$, which is the case for all observed double neutron star systems where a total mass has been accurately measured, we show that the ratio of the dimensionless tidal deformabilities satisfy $\Lambda_1/\Lambda_2\sim q^6$, where $q=M_2/M_1$ is the binary mass ratio; $\Lambda$ and $M$ are the dimensionless deformability and mass of each star, respectively. Moreover, they are bounded by $q^{n_-}\ge\Lambda_1/\Lambda_2\ge q^{n_{0+}+qn_{1+}}$, where $n_-<n_{0+}+qn_{1+}$; the parameters depend only on ${\cal M}$, which is accurately determined from the gravitational-wave signal. We also provide analytic expressions for the wider bounds that exist in the case of a strong phase transition. We argue that bounded ranges for $\Lambda_1/\Lambda_2$, tuned to ${\cal M}$, together with lower bounds to $\Lambda(M)$, will be more useful in gravitational waveform modeling than other suggested approaches. | Finite size effects in a binary neutron star merger are manifested, to lowest order, through the tidal deformabilities of the individual stars. The tidal effects are imprinted in the gravitational-wave signal through the binary tidal deformability~\citep{Flanagan08,Hinderer08} \begin{equation} \tilde\Lambda={16\over13}{(12q+1)\Lambda_1+(12+q)q^4\Lambda_2\over(1+q)^5}, \label{eq:tildelambda} \end{equation} where $q=M_2/M_1\le1$ is the binary mass ratio. The dimensionless deformability of each star is \begin{equation} \Lambda_{[1,2]}={2\over3}k_{2,[1,2]}\left({R_{[1,2]}c^2\over GM_{[1,2]}}\right)^5, \label{eq:lambda} \end{equation} where $k_2$ is the tidal Love number~\citep{Damour09,Flanagan08,Hinderer08}, which is the proportionality constant between an external tidal field and the quadrupole deformation of a star. $R_{[1,2]}$ and $M_{[1,2]}$ are the radii and masses of the binary components, respectively. $k_2$ can be readily determined from a first-order differential equation simultaneously integrated with the two usual TOV structural equations~\citep{Hinderer10,Postnikov10} and has values ranging from about 0.05 to 0.15 for neutron stars. For black holes, $k_2=0$. The tidal deformations of the neutron stars result in excess dissipation of orbital energy and speed up the final stages of the inspiral. Tidal deformations act oppositely to spin effects, which tend to be more important during earlier stages of the observed gravitational wave signal. The gravitational waves from the recently observed merger of two neutron stars, GW170817, were analyzed by the LIGO/VIRGO collaboration~\citep{Abbott17} (hereafter LVC), and subsequently reanalyzed by De et al. \cite{De18} (hereafter DFLB$^3$) and also the LIGO/VIRGO collaboration~\citep{Abbott18} (Hereafter LVC2). In the LVC analysis, the gravitational-wave signal was fitted to the Taylor F2 post-Newtonian aligned-spin model~\citep{Sathyaprakash91,Buonanno09,Arun09,Mikoczi05,Bohe13,Vines11} which has 13 parameters. 7 of those parameters are extrinsic, including the sky location, the source's distance, polarization angle and inclination, and the coalescence phase and time. The remaining 6 parameters are intrinsic, including the masses $M_1$ and $M_2$, dimensionless tidal deformabilities $\Lambda_{[1,2]}$, and the component's aligned spins $\chi_{[1,2]}=cJ_{[1,2]}/GM_{[1,2]}^2$, where $J$ is the angular momentum. The reanalysis of DFLB$^3$ differed from that of LVC chiefly in that electromagnetic observations were used to fix the source location and distance and in the adoption of the relation $\Lambda_1/\Lambda_2=q^6$, expressing the assumption that the two stars have a common equation of state (EOS). They justified this assumption using parameterized hadronic EOSs modeled using a fixed neutron star crust and three high-density polytropic segments whose parameters were restricted by causality and a minimum value of an assumed neutron star maximum mass. DFLB$^3$ also employed the causal lower limit to $\Lambda(M)$ in their analysis. In contrast, the analysis of LVC assumed uncorrelated priors for $\Lambda_1$ and $\Lambda_2$, thereby assuming that the two stars did not have the same equation of state, and did not consider causality-violating values of $\Lambda_1$ or $\Lambda_2$. DFLB$^3$ showed that models including correlations were favored by odds ratio $\gtrsim100$ over models using uncorrelated deformabilities, and, furthermore, that including deformability correlations reduced the 90\% confidence upper limit to the binary deformability by about 20\%. The latter result was confirmed by LVC2, who reanalyzed the GW170817 signal including deformability correlations using two different prescriptions. It is reasonable to assume that future investigations of neutron star mergers will treat $\Lambda_1$ and $\Lambda_2$ as correlated parameters, irrespective of which waveform model is used. The purposes of this paper are 1) to replace the approximate result $\Lambda_1/\Lambda_2=q^6$ with analytic bounds suitable for use in existing methods of fitting gravitational-wave signals of neutron star mergers, 2) to establish realistic lower limits to $\Lambda(M)$, 3) to compare our method with one proposed by Yagi and Yunes \cite{Yagi17}, and 4) to determine modifications to deformability correlations due to the possible existence of a strong first order phase transitions in the density range between the central densities of the two stars. In this case, the more massive star will be considered to be a {\it hybrid} star, in contrast to the lower mass star which we refer to as a {\it hadronic} star. This oversimplified notation harks back to the possibility of a hybrid hadronic-quark matter star in which the quark matter-hadronic matter interface has a surface tension too large to permit a smooth Gibbs phase transition. In the event of a strong first order phase transition, the more massive star can have a radius and tidal deformability much smaller than the lower mass star, even though their masses are nearly equal. This weakens the correlations otherwise evident between the tidal deformabilities and masses. In addition to bounds on the deformability ratio $\Lambda_1/\Lambda_2$, future analyses will benefit from the incorporation of absolute lower bounds to $\Lambda(M)$ available from consideration of the {\it maximally compact} EOS~\citep{Koranda97,Lattimer12}, which are limited by causality and the observed minimum value of the neutron star maximum mass. This EOS assumes that the matter pressure is essentially zero below a fiducial density $n_o$ that is a few times the nuclear saturation density, and that above this density the sound speed is equal to the speed of light. However, we also determine a more realistic and less extreme lower bound in which the pressure in the vicinity of the nuclear saturation density is instead limited from below by the unitary gas constraint thought to be applicable for neutron star matter~\citep{Tews17}. Upper bounds to $\Lambda(M)$ are available from nuclear theory and experiment, but are unfortunately model-dependent, and astrophysical observations also cannot yet provide accurate upper bounds. We will, however, explore the sensitivity of both lower and upper deformability bounds to assumptions concerning the minimum pressure of neutron star matter and also the minimum and maximum values assumed for the neutron star maximum mass. This paper is organized as follows: \S \ref{sec:likely} describes the most likely masses and spins for merging neutron star systems, and \S \ref{sec:tide} reviews how tidal deformabilites are defined and calculated. \S \ref{sec:par} outlines the parameterized equations of state used in this paper and the resulting tidal deformabilities and their bounds, while \S \ref{sec:bin} outlines results for the binary tidal deformabilities and their bounds. \S \ref{sec:hadron} establishes the correlations of tidal deformabilities with masses and compares our approach with other work. The lower bounds on deformabilities from causality are summarized in \S \ref{sec:caus}, and those from the unitary gas and neutron matter constraints are discussed in \S \ref{sec:unitary}. Deformability constraints for hybrid stars are established in \S \ref{sec:hybrid}. We summarize our conclusions in \S \ref{sec:conclusion}. | } In this paper, we have established upper and lower bounds for $\Lambda_2/\Lambda_1$ as functions of $q$ and ${\cal M}$, and minimum values of $\Lambda(M)$, that can be used to restrict the priors of deformabilities in analyses of gravitational-wave data from neutron star mergers. DFLB$^{3}$ has shown that taking these correlations and bounds into account significantly improves fits in the case of GW170817. Imposing correlations reduced the uncertainty range for $\tilde\Lambda$, lowering the 90\% credible upper limit by approximately 20\%. \begin{figure} \includegraphics[width=\linewidth,angle=180]{lambound-cropped.pdf} \caption{The variation of the upper and lower bounds to $\Lambda(M)$ for hadronic stars as the boundary densities $n_1$ and $n_2$ are changed. $M_{max}=2.0M_\odot$ is assumed.\label{fig:hbound}} \end{figure} The bounds we established for hadronic stars were based on a piecewise polytropic scheme with three segments and fixed boundary densities. We find our results with three segments to be relatively insensitive to reasonable variations of the boundary densities (Fig. \ref{fig:hbound}) $n_1$ and $n_2$. Varying the boundary densities produce variations of order $\pm5\%$ in the upper boundary and $\pm10\%$ in the lower boundary although, for a $1.4M_\odot$ star, the maximum value of $\Lambda$ is about 6 times the lowest value for $M_{max}=2M_\odot$. However, the variations produced by altering the number of polytropic segments can be more extreme. Adding polytropic segments allows for the possibility of one or more strong first-order phase transitions and so the upper and lower bounds to $\Lambda(M)$ can approach the results for the hybrid configurations in these cases. However, restricted to parameter ranges that approximate purely hadronic equations of state, varying the number of polytropic segments produce changes to $\Lambda(M)$ bounds similar to the changes induced by altering the boundary densities in the three-polytrope scheme shown in Fig. \ref{fig:hbound}. Modifying the piecewise polytrope scheme to smooth its behavior near the segment boundaries, as in the spectral decomposition method \cite{Lindblom18}, also has been shown to increase the accuracy in reproducing specific equations of state. Other high-density approximation methods have also been suggested, e.g., Ref.~\cite{Kurkela14}. However, such schemes inevitably reduce the allowed ranges of sampled pressure-density relations and therefore result in artificially smaller bounding ranges. It is important to emphasize that determining $\Lambda(M)$ bounds is dissociated from the question of a parameterized scheme's accuracy in reproducing $\Lambda(M)$ from a specific equation of state. Nevertheless, if one attempts to directly deduce the EOS itself from gravitational waveform modeling, as LVC2 has attempted, the accuracy of the high-density approximation scheme becomes an important consideration. | 18 | 8 | 1808.02858 |
1808 | 1808.02891_arXiv.txt | For the first time, we establish a gas phase abundance pattern calibration for quiescent galaxies using optical emission lines. Quiescent galaxies have warm ionized gas showing line ratios similar to low-ionization nuclear emission line regions (LINER). The ionization mechanism for the gas is still an unsettled puzzle. Despite the uncertainty in the ionization mechanism, we argue that we can still infer certain gas phase abundance pattern from first principles. We show that the relative trend in N/O abundance can still be reliably measured based on \niiw/\oiiw\ and a direct measurement of the electron temperature. We construct a composite direct temperature tracer that is independent of extinction correction, by combining \oiiw/\oiitw\ and \siiw/\siitw\ and canceling out the effect of extinction, as these involve the easiest-to-detect auroral lines in quiescent galaxies. With theoretical modeling, we establish the calibration for N/O based on \nii/\oii\ and a temperature tracer. We apply this technique to quiescent galaxies in the nearby Universe and find they span a range of 0.35 dex in N/O ratio from 17-percentile to 83-percentiles of the whole distribution. These measurements can shed light on the chemical enrichment history of the warm ionized gas in quiescent galaxies. | Galaxies that are not forming stars can be referred to as quiescent galaxies. In the past, quiescent galaxies are often thought to contain no neutral gas. But this view has been challenged by recent observations of molecular and atomic gas in these galaxies\citep[e.g.][]{DavisT11,Serra12,Young14}. Optical spectroscopy observations also revealed weak optical emission lines in these galaxies indicating warm ionized gas with a temperature around $10^4$K and dust\citep{Phillips86,Kim89,Buson93,Goudfrooij94,Macchetto96,Zeilinger96,Lauer05,Sarzi06,DavisT11, Singh13, Belfiore16, Gomes16}. The relative intensity ratios among the optical emission lines indicate their production must be different from the photoionization by young hot stars in star-forming galaxies. However, we do not yet know for sure what physical mechanism produced these warm ionized gas. Despite this uncertainty, there can be abundance pattern variations we can measure based on direct temperature measurements. Because the abundance pattern of the interstellar medium provides the fossil record of the chemical enrichment history, measuring the elemental abundance of the gas will provide important insights into the star formation history of these galaxies and the baryonic flows in and out of them. Gas phase abundances can be estimated using many different wavebands. Using X-ray observations, one could measure the iron abundance in the hot gas in massive ellipticals and galaxy clusters \citep[e.g.][]{EdgeS91, AllenF98}. However, the lighter elements are much more difficult to measure in X-rays. The lighter elements are much more easily detectable with optical emission lines, which probe the warm ionized gas component. With optical emission lines, gas metallicity are usually measured only for star-forming galaxies as we understand well the ionization mechanism in star-forming regions. For quiescent galaxies, there have been very few attempts before. \cite{AtheyB09} applied the R23 calibration derived from star-forming region models on early-type galaxies. They assumed that the gas is photoionized by post-Asymptotic Giant Branch stars (post-AGB stars) and neglected the difference in the ionizing spectra between post-AGB stars and that of massive OB stars. This can lead to very large systematic uncertainty in the resulting oxygen abundance, and can also change the relative abundance difference between galaxies. \cite{Storchi-Bergmann98} did similar analysis for LINERs but that study suffers from the same issue. The ionization mechanism for the warm ionized gas in quiescent galaxies has been hotly debated for more than 20 years. Several ionization mechanisms have been proposed to explain the line ratios observed in them. These include photoionization by an active galactic nucleus \citep{FerlandN83,HalpernS83,GrovesDS04II}, photoionization by post-AGB stars \citep{Binette94, Stasinska08}, photoionization by the hot X-ray emitting gas \citep{VoitD90, DonahueV91}, collisional ionization by fast shocks \citep{DopitaS95}, and heat exchange through conduction or turbulent mixing layers between hot and cold gas \citep{SparksMG89,Slavin93}. Among these, photoionization by AGN has been largely ruled out as the dominant mechanism in the great majority of quiescent galaxies \citep{Sarzi10,YanB12,Singh13,Belfiore16,Gomes16}. However, the jury is still out on which of the other mechanisms is the dominant source. Photoionization by post-AGB stars is often considered to be the most likely source, but this is not confirmed. A recent investigation by \cite{Yan18a} found that neither photoionization nor shocks could explain all the line ratios measured in the stacked spectra of these quiescent galaxies. Despite this uncertainty in the ionization mechanism, we argue in this paper that it is still possible to measure the abundance ratios among a few elements. For example, we can measure N/O abundance ratio as long as we can get a measurement of \niiw/\oiiw\ and a measurement of electron temperature. This is because, regardless of the ionization mechanism, the \niiw\ and \oiiw\ lines are always collisionally-excited and the ratio between them only depends on N$^+$/O$^+$ and the temperature. Nitrogen and Oxygen also have very similar structures in their ionization potentials: the energies to ionize neutral N and O to N$^+$ and O$^+$ differ by only $0.9$eV, while the energies to ionize these to the next level differ by $5.5$eV. This means that, as long as the energy distribution of the ionizing photons or particles is reasonably smooth, Nitrogen and Oxygen would have nearly the same fraction of atoms in a singly ionized state and those ions would be found in the same spatial region, regardless of the ionization mechanism. This assumption should hold for all the ionization mechanisms mentioned above. Therefore, we can establish a calibration for N/O abundance ratio in quiescent galaxies, although we have not fully settled the ionization mechanism for the gas. We use photoionization models to derive the calibration, while knowing the calibration is not just limited to photoionization models. Different ionization mechanism may cause small systematic shifts in the absolute N/O abundance ratios due to inaccuracies in the above assumptions. But the relative trend should remain robust. We then apply this to measure the N/O abundance in quiescent galaxies in the nearby Universe. Our approach is similar to that of the metallicity calibrations for star-forming galaxies, but we improve on it in three important aspects. First, we do not assume a fixed relationship between N/O and O/H in our theoretical modeling. Instead, we establish a grid of models sampling a wide range of N/O vs. O/H patterns. Second, we construct a direct electron temperature tracer that is independent of extinction correction. This is based on two direct temperature tracers which have been rarely used due to their sensitivity to extinction. We combine them together to cancel out the extinction effect and form an extinction-insensitive temperature tracer. Third, the calibration we derive will link line ratio measurements directly with abundance ratios, without going through an intermediate step of deriving temperatures and ionization correction factors. The latter has been the approach adopted in many $T_e$-based abundance calibrations. Our approach will reduce the number of necessary assumptions and the systematic errors associated with those assumptions. In this paper, we first describe the setup of the simulations(\S\ref{sec:simulation}), then establish the N/O abundance calibration(\S\ref{sec:calibration}), and finally apply the calibration on SDSS galaxies to measure their N/O abundance ratio (\S\ref{sec:data}). | Gas phase abundance patterns in quiescent galaxies are difficult to explore in the optical wavelengths due to the uncertainty of the ionization mechanism. Despite this uncertainty, we believe one can make progress on some abundance ratios, such as N/O, based on simple physical arguments that apply to multiple ionization mechanisms. In this paper, we have developed a simple N/O abundance ratio calibration for quiescent galaxies. It is a direct T$_e$-based method as it requires the detection of temperature-sensitive auroral lines. We have devised a new extinction-insensitive temperature tracer based on auroral lines from O$^+$ and S$^+$. We also explored a wide range of parameter space in N/O vs. O/H to prove that a reliable measurement of N/O can be made independent of the assumed N/O vs. O/H relationship. We have applied our calibration to quiescent galaxies in SDSS and found they have a 0.35 dex spread in N/O ratio between 17- and 83-percentiles. This will help shed light on the chemical enrichment history of the warm ionized gas in these galaxies. One may want to take this method further and measure O/H abundance in these quiescent galaxies. We refrain from doing this because that derivation would depend strongly on the assumption of the ionization mechanism. We warn the reader from taking this method too far before confirming the ionization mechanism. | 18 | 8 | 1808.02891 |
1808 | 1808.00636_arXiv.txt | The spectral evolution and spectral lag behavior of 92 bright pulses from 84 gamma-ray bursts (GRBs) observed by the Fermi GBM telescope are studied. These pulses can be classified into hard-to-soft pulses (H2S, 64/92), H2S-dominated-tracking pulses (21/92), and other tracking pulses (7/92). We focus on the relationship between spectral evolution and spectral lags of H2S and H2S-dominated-tracking pulses. The main trend of spectral evolution (lag behavior) is estimated with $\log E_p\propto k_E\log(t+t_0)$ (${\hat{\tau}} \propto k_{\hat{\tau}}\log E$), where $E_p$ is the peak photon energy in the radiation spectrum, $t+t_0$ is the observer time relative to the beginning of pulse $-t_0$, and ${\hat{\tau}}$ is the spectral lag of photons with energy $E$ with respect to the energy band $8$-$25$~keV. For H2S and H2S-dominated-tracking pulses, a weak correlation between $k_{{\hat{\tau}}}/W$ and $k_E$ is found, where $W$ is the pulse width. We also study the spectral lag behavior with peak time $t_{\rm p_E}$ of pulses for 30 well-shaped pulses and estimate the main trend of the spectral lag behavior with $\log t_{\rm p_E}\propto k_{t_p}\log E$. It is found that $k_{t_p}$ is correlated with $k_E$. We perform simulations under a phenomenological model of spectral evolution, and find that these correlations are reproduced. We then conclude that spectral lags are closely related to spectral evolution within the pulse. The most natural explanation of these observations is that the emission is from the electrons in the same fluid unit at an emission site moving away from the central engine, as expected in the models invoking magnetic dissipation in a moderately-high-$\sigma$ outflow. | \label{sec:intro} It was theoretically speculated and observationally confirmed that long gamma-ray bursts (GRBs) are from the collapse of massive stars and short GRBs are from mergers of compact binaries (e.g., \citealp{Colgate_SA-1974}; \citealp{Paczynski_B-1986}; \citealp{Eichler_D-1989-Livio_M}; \citealp{Narayan_R-1992-Paczynski_B}; \citealp{Woosley_SE-1993}; \citealp{Woosley_SE-2006-Bloom_JS}; \citealp{Kumar_P-2015-Zhang_B}). A spectral lag, defined as the time delay of high-energy photons with respect to low-energy photons, is commonly observed in long GRBs (\citealp{Norris_JP-1986-Share_GH}; \citealp{Cheng_LX-1995-Ma_YQ}; \citealp{Band_DL-1997}; \citealp{Norris_JP-2000-Marani_GF}), but is not in short ones (\citealp{Yi_T-2006-Liang_E}; \citealp{Norris_JP-2006-Bonnell_JT}). An extensive analysis of the data from the Burst And Transient Source Experiment (BATSE) also shows that long wide-pulse GRBs tend to have long spectral lags (\citealp{Norris_JP-2005-Bonnell_JT}). The distinguished features of spectral lags in these two types of GRBs are proposed to be a phenomenological indicator in GRB classification (\citealp{Yi_T-2006-Liang_E}; \citealp{Norris_JP-2006-Bonnell_JT}; \citealp{gehrels06}; \citealp{McBreen_S-2008-Foley_S}; \citealp{zhangbb09}). An anti-correlation between the spectral lag and peak luminosity is found in a few redshift-known BATSE GRBs (\citealp{Norris_JP-2000-Marani_GF}). The low-luminosity GRB~060218 detected with the Burst Alert Telescope (BAT) onboard Swift satellite is consistent with this correlation (\citealp{Liang_EW-2006-Zhang_BB}). The precursor and main prompt emission in GRB 061121 show different spectral lags, but both are consistent with the above correlation (\citealp{Page_KL-2007-Willingale_R}). The above correlation also holds in \emph{Swift}/BAT GRB sample and even in the X-ray flares observed with the X-Ray Telescope onboard {\em Swift}, albeit with large scatter (\citealp{Ukwatta_TN-2010-Stamatikos_M}; \citealp{Margutti_R-2010-Guidorzi_C}; \citealp{Sultana_J-2012-Kazanas_D}; \citealp{Sonbas_E-2013-MacLachlan_GA}; \citealp{Bernardini_MG-2015-Ghirlanda_G}). The relation between spectral lags and peak luminosity was proposed to be a distance indicator of GRBs (\citealp{Norris_JP-2000-Marani_GF}; \citealp{Schaefer_BE-2007}) for the purpose of using GRBs to constrain cosmological parameters. GRB light curves are usually composed of overlapping pulses (\citealp{Norris_JP-1996-Nemiroff_RJ}; \citealp{Hu_YD-2014-Liang_EW}). It is generally speculated that each pulse is related to one individual radiation episode, which represents the fundamental unit of the GRB temporal profile (see, e.g., \citealp{Fishman_GJ-1994-Bhat_PN}). The spectral lags measured in different emission episodes could be different (e.g., \citealp{Page_KL-2007-Willingale_R}). However, the measured spectral lag is likely dominated by the wide pulses or the brightest pulse in a light curve. As a result, \cite{Hakkila_J-2008-Giblin_TW} suggested that the spectral lag is better defined using individual pulses rather than the whole burst light curve profile. Significant spectral evolution is also a common feature of GRB pulses (e.g., \citealp{Liang_E-1996-Kargatis_V}; \citealp{Preece_RD-2000-Briggs_MS}; \citealp{Lu_RJ-2010-Hou_SJ,Lu_RJ-2012-Wei_JJ}; \citealp{Racz_II-2018-Balazs_LG}). Two evolutionary patterns have been observed. One is the so-called hard-to-soft (hereafter H2S) pattern, which can be defined as the peak photon energy ($E_{p}$) in $\nu$-$\nu f_\nu$ spectrum decreasing monotonically within an individual pulse. The second is the so-called intensity tracking pattern, which is defined by $E_{p}$ tracking the flux (\citealp{Wheaton_WA-1973-Ulmer_MP}; \citealp{Norris_JP-1986-Share_GH}; \citealp{Golenetskii_SV-1983-Mazets_EP}; \citealp{Lu_RJ-2012-Wei_JJ}). It was suggested that the spectral lag behavior of GRB prompt emission may be related to spectral evolution (\citealp{Ukwatta_TN-2012-Dhuga_KS}; \citealp{uhm16}; \citealp{Preece2016}). Some authors also pointed out that the spectral evolution is related to the pulse profile (Hakkila \& Preece 2011; Hakkila et al. 2015, 2018). For example, the hard-to-soft spectral evolution may primarily occur in hard and/or asymmetric pulses, and the intensity-tracking spectral evolution may be more prevalent in soft and/or symmetric pulses (Hakkila et al. 2015, 2018). In addition, Hakkila \& Preece (2014) found that the residuals of single pulse fits to GRB lightcurves leave behind a mysterious triple-peaked structure. They argued that the interplay between such a triple-peaked structure and pulse fits may complicate spectral evolution in pulses. In this paper, we explore the relation between the spectral lags and spectral evolution. Gamma-Ray Burst Monitor (GBM) onboard the Fermi satellite has established a large GRB sample. We present a comprehensive analysis of the \emph{Fermi} GRB data and report our results in a series of papers. We revealed the spectral components and their temporal evolution of the \emph{Fermi} GRBs in the first two papers of this series (\citealp{zhangbb11}; \citealp{Lu_RJ-2012-Wei_JJ}) and studied energy dependence of the burst duration in the third paper (\citealp{Qin_Y-2013-Liang_EW}). This paper is dedicated to investigating the spectral lag and its relation to spectral evolution within bright GRB pulses. We describe our sample selection and data reduction in Section~2. The relations between the spectral evolution and the spectral lag behavior are studied in Section~3. Based on the results in Section~2, we explore the possible origin of spectral lags by performing simulations within the framework of a phenomenological model in Section~4. We summarize our results in Section~5. Throughout of this paper, a flat $\Lambda$CDM cosmology with the parameters $H_{0}=71 \rm km \cdot s^{-1}\cdot Mpc^{-1}$, $\Omega_{\rm M}=0.3$, and $\Omega_{\Lambda}=0.7$ is adopted. The quoted uncertainties are at $1\sigma$ confidence level. | This work studies the spectral evolution and spectral lag behavior in 92 bright pulses from 84 GRBs observed by the Fermi GBM telescope. We focus on the relation of spectral evolution and the spectral lag behavior in H2S (64/92) and H2S-dominated-tracking (21/92) pulses. It is found that the spectral lag (${\hat{\tau}}$) is usually photon-energy-dependent. In the H2S and H2S-dominated-tracking pulses, the spectral lag increases with increasing photon energy $E$ (given the same reference band in 8-25 keV). In addition, the dependence of spectral lag on $E$ may be different from pulse to pulse, even for those in the same GRB. We then adopt the slope $k_{\hat{\tau}}$ in the relation of ${\hat{\tau}} \propto k_{\hat{\tau}}\log E$ to describe the behavior of $E$-dependent ${\hat{\tau}}$ for different pulses. The main trend of $E_p$ evolution is approximated by performing a log-linear fit to the $E_p$-$t$ relation with $\log E_p(t)\propto k_E\log(t+t_0)$ and the $k_E$ is used to describe the behavior of spectral evolution. For H2S and H2S-dominated-tracking pulses, a weak relation of $k_{{\hat{\tau}}}/W$ and $k_E$ is found, where $W$ is the full width at half maximum for the pulses observed at $8$-$10^3$~keV energy band. For further studying the relation between spectral evolution and spectral lag, we also investigate the evolution of peak time $t_{\rm p_E}$ in different energy bands for 30 well shaped pulses. A log-linear relation of $t_{\rm p_E}$-$E$ is found, i.e. $\log t_{\rm p_E}\propto k_{t_p}\log E$, in H2S and H2S-dominated-tracking pulses. In addition, a weak relation between $k_{t_p}$ and $k_E$ is also obtained. The relations of $k_{{\hat{\tau}}}$-$k_E$ and $k_{t_p}$-$k_E$ together with the spectral evolution pattern are reproduced in our simulations within the framework of a phenomenological model that invokes the evolution of a Band-function spectrum with time. Based on these results, we conclude that spectral lag is related to the spectral evolution in GRBs. The discovery reported in this paper sheds light on the GRB prompt emission mechanism that is highly debated. The fact that the spectrum evolves uniformly within each pulse (which has a typical duration of seconds) suggests that pulses are fundamental units of GRB radiation. Within the framework of GRB models, there are two scenarios to interpret these broad pulses. The first scenario is that the pulse shape is defined by the central engine activity history, so that each emission epoch in the light curve corresponds to emission from different groups of electrons when they reach a certain radius (e.g., photosphere or internal shocks). The second scenario is that the broad pulse is emission from electrons within the same fluid unit, and different emission episodes correspond to the same fluid unit emitting at different locations as it streams outwards. This second possibility corresponds to an emission radius $R_{\rm GRB,pulse} \sim \Gamma^2 c t_{\rm pulse} \sim 10^{15} \ {\rm cm} (\Gamma/100)^2 (t_{\rm pulse}/ 3 \ {\rm s})$. Any model that invokes an emission radius smaller than this radius, e.g., the photosphere model (unless $\Gamma$ is extremely small) and the internal shock model (unless the variability time $\Delta t$ is of the order of pulse duration), belong to the first scenario. Our data and simulations suggest that spectral lags are closely related to the spectral evolution of the emission spectrum. This is much easier to realize if emission comes from electrons within the same fluid unit, so that emission properties can evolve continuously as the fluid unit moves in space (so that the magnetic field strength, bulk Lorentz factor, and probably characteristic electron Lorentz factors) can evolve continuously (see the generic physical model discussed by \citealp{uhm16} and \citealp{uhm18}). Such a scenario naturally produces asymmetric pulse profiles and H2S evolution. Under certain conditions, one can also reproduce tracking pulses (most are H2S-dominated)\footnote{ The mysterious triple-peaked structure claimed by \cite{Hakkila2014} is, however, difficult to interpret within the framework of any physically oriented models.}. It is consistent with dissipation of magnetic energy at a radius of the order of $R_{\rm GRB,pulse}$, as has been invoked in the ICMART model of GRB prompt emission (\citealp{zhangyan11}; \citealp{lazarian18}). Within this picture, the smaller variability timescale overlapping on the broad pulses are produced by the local Lorentz-boosted regions due to magnetic reconnection within a moderately-high-$\sigma$ bulk flow (\citealp{zhangyan11}; \citealp{zhangzhang14}; \citealp{deng15}). Within the framework of the first scenario, on the other hand, the requirement to have emission spectrum evolves corporately as a function of time is much more demanding, since electrons are from different fluid units. Indeed within the framework of the photosphere model, it has been shown that it is quite challenging to produce H2S evolution pulses (\citealp{deng14}). The small-radii internal shock model (which is needed to interpret the rapid variability in the light curves) also faces the similar problem, since it requires that the electron Lorentz factors, magnetic fields, and bulk Lorentz factors in different internal shocks would behave corporately with conspiracy to give rise to the observed spectral evolution. | 18 | 8 | 1808.00636 |
1808 | 1808.10817_arXiv.txt | s{NIKA2 is a dual-band millimetric camera of thousands of Kinetic Inductance Detectors (KID) installed at the IRAM 30-meter telescope in the Spanish Sierra Nevada. The instrument commissioning was completed in September 2017, and NIKA2 is now open to the scientific community and will operate for the next decade. NIKA2 has well-adapted instrumental design and performance to produce high-resolution maps of the thermal Sunyaev-Zel'dovich (SZ) effect toward intermediate and high redshift galaxy clusters. Moreover, it benefits from a guaranteed time large program dedicated to mapping a representative sample of galaxy clusters via SZ and that includes X-ray follow-ups. The main expected outputs of the SZ large program are the constraints on the redshift evolution of the pressure profile and the mass-observable relation. The first SZ mapping of a galaxy cluster with NIKA2 was produced, as part of the SZ large program. We found a sizable impact of the intracluster medium dynamics on the integrated SZ observables. This shows NIKA2 capabilities for the precise characterisation of the mass-observable relation that is required for accurate cosmology with galaxy clusters.} \clearpage | 18 | 8 | 1808.10817 |
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1808 | 1808.00510_arXiv.txt | We study the spectroscopic evolution of superluminous supernovae (SLSNe) later than 100 days after maximum light. We present new data for Gaia16apd and SN\,2017egm, and analyse these with a larger sample comprising 41 spectra of 12 events. The spectra become nebular within 2-4 $e$-folding times after light curve peak, with the rate of spectroscopic evolution correlated to the light curve timescale. Emission lines are identified with well-known transitions of oxygen, calcium, magnesium, sodium and iron. SLSNe are differentiated from other Type Ic SNe by a prominent \oi\ line and higher-ionisation states of oxygen. The iron-dominated region around 5000\,\AA\ is more similar to broad-lined SNe Ic than to normal SNe Ic. Principal Component Analysis shows that 5 `eigenspectra' capture $\gtrsim70$\% of the variance, while a clustering analysis shows no clear evidence for multiple SLSN sub-classes. Line velocities are 5000--8000\,\kms, and show stratification of the ejecta. \oi\ likely arises in a dense inner region that also produces calcium emission, while \o\ comes from further out until 300--400 days. The luminosities of \oi\ and \ion{Ca}{2} suggest significant clumping, in agreement with previous studies. Ratios of \ca/\o\ favour progenitors with relatively massive helium cores, likely $\gtrsim 6$\,\M, though more modelling is required here. SLSNe with broad light curves show the strongest \o, suggesting larger ejecta masses. We show how the inferred velocity, density and ionisation structure point to a central power source. | \label{s:intro} In recent years, much progress has been made in characterising the new population of hydrogen-poor superluminous supernovae (SLSNe). These events first started to appear in wide-field, untargeted transient surveys \citep{qui2011,chom2011,gal2012}, and now comprise a few percent of the supernovae classified each year. SLSNe have a median luminosity $M \approx -21$\,mag \citep{nic2015b,lun2018,dec2018}, making them up to two orders of magnitude brighter than typical supernovae (SNe), and sparking intense interest in the unexpectedly diverse outcomes of massive star deaths. SLSNe are now generally classified spectroscopically rather than photometrically. Their unique early-time spectra show a series of broad \ion{O}{2} absorption lines superposed on a blue continuum, indicating hot, ionized ejecta. As they expand and cool, the spectra evolve to resemble those of lower luminosity Type Ic SNe \citep{pas2010,ins2013}, though there may be subtle differences \citep{Liu&Modjaz2017,qui2018}. Thus SLSNe are best described as Type Ic SNe (explosions of stripped massive stars lacking hydrogen and helium) that manage to stay hot ($T \gtrsim 10,000$\,K) over several weeks or months, allowing them to attain higher luminosities. The favoured source of additional heating is a central engine, such as the spin-down of a rapidly rotating magnetar \citep{kas2010,woo2010}. It has been shown that this model can reproduce both the light curves \citep{ins2013,cha2013,nic2017c} and early spectra \citep{des2012,how2013,maz2016} of SLSNe. Maximum-light observations, however, only probe the outer layers of the ejecta, because of a large optical depth to the centre. Over time, recombination reduces the optical depth, and by $t_{\rm neb} \sim 360\,{\rm d} (M_{\rm ej}/10$\,\M$)^{1/2} (v/10^4$\,\kms$)^{-1}$ after the explosion, where \Mej\ is the ejecta mass and $v$ is the expansion velocity, the ejecta become largely transparent \citep[see review by][]{jer2017b}. Once this so-called nebular phase is reached, it is possible to directly probe with spectroscopy the conditions at the centre of the explosion, to constrain the composition and distribution of material and to search for any hydrodynamic signatures of the explosion mechanism. However, such observations are challenging because the SN will have faded substantially over the time $t_{\rm neb}$, and hence nebular spectroscopy is only currently possible for SLSNe at $z \lesssim 0.2$. \citet{nic2016c} and \citet{jer2017a}, following earlier work by \citet{mil2013}, showed that the few existing nebular spectra of SLSNe resemble those of broad-lined Type Ic SNe -- thought to be engine-driven explosions and sometimes accompanied by long gamma-ray bursts (LGRBs). This was interpreted as evidence for a similar internal structure for SLSNe and LGRB SNe, suggesting that they may arise from similar progenitors. They also inferred large ejected masses of $\gtrsim 10$\,\M\ for some SLSNe. Since those initial nebular observations, the available sample of SLSNe with nebular phase observations has increased as a number of recent nearby SLSNe have evolved to sufficiently late times. For example, \citet{qui2018} recently published a large spectroscopic sample of SLSNe from the Palomar Transient Factory (PTF), including many late-phase spectra. In this paper, we undertake a systematic observational study of SLSN nebular spectra. We present new data for Gaia16apd/SN\,2016eay ($z=0.1013$) and SN\,2017egm ($z=0.0307$), two nearby SLSNe that have been well-studied at earlier phases \citep{yan2016,nic2017,kan2016b,nic2017d,bos2018}, and combine this with all available published spectra of SLSNe obtained more than 100 days after maximum light. We describe our observations and the processing of new and archival data in Section \ref{s:data}. In Section \ref{s:obs}, we present the spectral sequence and mean population properties, including line identifications and comparisons to Type Ic SNe. We then apply machine learning techniques to characterise the diversity of SLSNe in Section \ref{s:mach}. The line profiles are used to investigate the distribution of material in Section \ref{s:profiles}, and their luminosities and ratios to infer ejecta conditions in Section \ref{s:lums}. We discuss the implications of our findings in the context of SLSN models in Section \ref{s:diss} and summarise our conclusions in Section \ref{s:conc}. | \label{s:conc} We have conducted a systematic study of the observed properties of SLSN spectra as they evolve through the nebular phase. Our sample comprised 41 spectra of 12 SLSNe, including both fast and slow evolving events. After applying a consistent interpolation and smoothing procedure to all spectra, and normalising the observed phase by the different decay rates of the SLSN light curves, we found that all events could be reasonably well described in terms of a single spectral sequence, ordered by this normalised phase. Our GMOS spectra of Gaia16apd (399 days after peak) and SN\,2017egm (353 days after peak) are among the latest obtained for SLSNe, especially when considering the relatively fast light curve timescales of these events. We analysed these spectra in terms of their statistical properties and the velocity and luminosity evolution of specific lines. For convenience, we summarise our main conclusions here: \begin{itemize} \item SLSN spectra become dominated by nebular features within 2--4 $e$-folding times after their light curve peaks. This provides a means early in the evolution of an individual event to plan the optimal time for nebular-phase follow-up. \item The main emission lines are easily identified with well-known transitions of oxygen, calcium, magnesium, sodium and iron -- the same species typically seen in normal and broad-lined SNe Ic. \item \o\ is initially weaker than \ca, and takes longer to develop a Gaussian line profile, but usually becomes the strongest optical line at later times. The ratio of the 6300\,\AA/6364\,\AA\ components indicates that this line is optically thin during most of the nebular phase. \item Compared to SNe Ic, SLSNe are differentiated by a prominent \oi\ recombination line, often along with [\ion{O}{2}] and [\ion{O}{3}], indicating higher ionisation, and the presence of oxygen in regions with significant variation in electron density. \item SLSNe also show elevated flux compared to normal SNe Ic, on average, over the iron-dominated part of the spectrum between 4000-5500\,\AA. However, the SLSN iron region is similar to some broad-lined SNe Ic such as SN\,1998bw. \item Principal component analysis showed that $\gtrsim 70$\% of the variance in SLSN spectra could be attributed to 5 eigenspectra, corresponding roughly to emission from \ion{Fe}{2}, \ion{Ca}{2} NIR triplet, \o, \ca, and \oi, in order of decreasing variance. \item We find no compelling evidence for clustering of SLSNe into sub-populations based on their nebular spectra. A K means clustering analysis with two assumed clusters separates the spectra as much by phase relative to explosion as by the actual SLSNe to which these spectra belong. \item Most SLSNe show no strong asymmetry in their \o, \ca, or \mg\ line profiles; LSQ14an is a possible exception to this. However, at least half of the SLSNe in our sample exhibit an excess on the red side of \oi, likely attributable to \ion{Mg}{2}\,$\lambda$7877,7896. \item The ejecta structure and composition inferred from the widths of the strongest lines is consistent with explosion models of massive stars, with calcium towards the centre and magnesium further out, and oxygen spanning a wide range of ejecta zones. \o\ arises from faster (further out) regions of the ejecta than \oi\ for several hundreds of days. \item The luminosity evolution in all lines decreases with time. After normalising the spectral phase by the light curve decline timescale, all SLSNe show virtually the same decline rate in their line luminosities, indicating that the nebular lines are reprocessing the same power source responsible for the earlier luminosity evolution of SLSNe. The ratios between many lines look similar to SNe Ic, provided one takes into account the relatively slower light curve evolution of SLSNe. \item The \ca/\o\ ratio matches some models with masses of $\sim 3.5-5.9$\,\M, but the measured ratio may be contaminated by [\ion{O}{2}]\,$\lambda$7320,7330 emission, such that the true core mass is likely even higher in some events. However, other events appear to be indicative of a population extending to lower masses, and selection effects may account for a large number of high-mass events in the nebular sample. \item The \oi\ line and the ratio between calcium lines indicates a region of high electron density with low filling factor at low velocity coordinate. We suggest that this could be due to the hydrodynamic impact of a central engine. \item The ionisation structure shows no evidence of changing over several hundred days. This is expected in the engine-powered models recently calculated by \citet{margalit2018}. In some SLSNe, oxygen may be ionised in a large fraction of the ejecta throughout the nebular phase -- these events are SN\,2017egm, LSQ14an and PTF10hgi. \item The radii reached by the ejecta on the timescales of these data are comparable to the radii where some SLSNe encounter hydrogen-rich material \citep{yan2015,yan2017}. That our spectra show no sign of hydrogen indicates a diversity in when the hydrogen layer is lost by the progenitors. \end{itemize} There remain important outstanding questions, including an accurate calibration of the ejecta masses for better comparison with light curve models, and determining the processes in the pre-supernova stellar evolution that expel the stellar envelope, and likely allow the formation of a central engine. It also remains to be demonstrated whether our suggested ejecta distribution is consistent with the observed smooth line profiles, and if \ca\ emission from deep-lying zones can escape the ejecta early in the transition to the nebular phase. Thus, a full understanding of SLSN progenitors will require more detailed modelling of their late-time spectra. Our analysis here provides an observationally-motivated starting point for exploring the model parameter space. \software{scikit-learn \citep{ped2011}, scipy \citep{scipy}, matplotlib \citep{matplotlib}, PyRAF, IRAF \citep{iraf}} | 18 | 8 | 1808.00510 |
1808 | 1808.09431_arXiv.txt | Both the interplanetary space and the Earth magnetosphere are populated by low energy ($\leq300$ keV) protons that are potentially able to scatter on the reflecting surface of Wolter-I optics of X-ray focusing telescopes and reach the focal plane. This phenomenon, depending on the X-ray instrumentation, can dramatically increase the background level, reducing the sensitivity or, in the most extreme cases, compromising the observation itself. The use of a magnetic diverter, deflecting protons away from the field of view, requires a detailed characterization of their angular and energy distribution when exiting the mirror. We present the first end-to-end Geant4 simulation of proton scattering by X-ray optics and the consequent interaction with the diverter field and the X-ray detector assembly, selecting the ATHENA Wide Field Imager as a case study for the evaluation of the residual soft proton induced background. We obtain that, in absence of a magnetic diverter, protons are indeed funneled towards the focal plane, with a focused Non X-ray Background well above the level required by ATHENA science objectives ($5\times10^{-4}$ counts cm$^{-2}$ s$^{-1}$ keV$^{-1}$), for all the plasma regimes encountered in both L1 and L2 orbits. These results set the proton diverter as a mandatory shielding system on board the ATHENA mission and all high throughput X-ray telescopes operating in the interplanetary space. For a magnetic field computed to deflect 99\% of the protons that would otherwise reach the WFI, Geant4 simulations show that this configuration, in the assumption of a uniform field, would efficiently shield the focal plane, yielding a residual background level of the order or below the requirement. | Low energy protons ($\leq300$ keV, so-called soft protons), populating the interplanetary space and the Earth magnetosphere, can enter the field of view of X-ray focusing telescopes and then be funneled towards the focal plane by scattering at grazing angles with the mirror surface. This phenomenon was discovered after the damaging of the Chandra/ACIS front-illuminated CCDs in 1999 during its first passages through the radiation belt \citep{2000SPIE.4140...99O}. The damage was soon minimized by switching off the CCDs and moving them from the focal position \citep{2007SPIE.6686E..03O}. Blocking filters protect XMM-Newton focal plane below an altitude of 40000 km, but above this limit soft protons induce sudden flares in the background count rate of the EPIC instruments. These events last from hundreds of seconds to hours and can hardly be disentangled from X-ray photons, causing the loss of large amounts (30-40\%) of observing time \citep{2017ExA....44..297M}. While telescopes operating in low Earth orbit are shielded by the Earth geomagnetic cut-off, the performance of future X-ray focusing telescopes orbiting in the interplanetary space can potentially suffer from soft proton induced background events. Examples of such missions are the ESA next large class ATHENA \citep{2013arXiv1306.2307N}, to be launched in 2030, or the eROSITA X-ray telescope on board the Russian/German Spectrum Roentgen Gamma Mission \citep{2014SPIE.9144E..1TP}, to be launched in 2019. \\ The large effective area (1.4 m$^{2}$ at 1 keV) makes the minimization of soft proton contamination a key challenge for the fulfillment of ATHENA's science objectives. A possible shielding solution is placing an array of magnets (a magnetic diverter) between the optics and the focal plane, able to deflect charged particles away from the instruments field of view. X-ray telescopes on board Chandra, XMM-Newton, and Swift are already equipped with diverters deflecting the electrons populating the radiation belts (see e.g \citet{wil00}). A proton diverter however, because of the $\sim2000$ times higher mass of the particle with respect to electrons, imposes a dedicated trade-off among the required magnetic field, the mass budget, and the impact on surrounding instruments. \\ The ATHENA Wide Field Imager (WFI) aims, among the many scientific objectives \citep{2016SPIE.9905E..2BR}, to perform X-ray surveys of the high-z sky, populated by faint point sources, and to map the diffuse and faint thermal emission in clusters of galaxies. A low instrumental background is mandatory for the achievement of those science objectives \citep{wfi_bkg2018}. Because of this requirement and its large field of view ($40'\times40'$), the WFI is the best case study for the evaluation of the soft proton induced X-ray background and the shielding efficiency of a magnetic diverter placed in front of it. For the first time we present an end-to-end simulation of the soft proton induced background, including (i) the collection of all plasma regimes encountered in L2 (Sec. \ref{sec:l2}), (ii) the interaction of protons with the mirror (Sec. \ref{sec:sca}), (iii) the consequent interaction with optical blocking filter in the field of view, the surrounding structure, and the detector itself (Sec. \ref{sec:wfi}) and (iv) the impact of a magnetic diverter in deflecting proton tracks from the WFI (Sec. \ref{sec:div}). The evaluation of the ATHENA diverter efficiency is one of the products of the ESA AREMBES (ATHENA Radiation Environment Models and X-ray Background Effects Simulators) project\footnote{http://space-env.esa.int/index.php/news-reader/items/AREMBES.html}, that will deliver to the science community a full modeling of the ATHENA space radiation environment and a Geant4-based framework, including the full ATHENA mass model, for the simulation of the Non X-ray Background (NXB). | \label{sec:con} The efficiency of the proton diverter in minimizing the focused NXB depends on the maximum proton energy to deflect and on the proton incident angle, not on the intensity of the beam. Comparing the impact of different plasma regimes, a harder X-ray spectrum, with more protons at higher energies despite the lower energy integrated flux, could reduce the diverter efficiency. Within the framework explored with our simulations we find that the efficiency of the diverter could also decrease if we place additional stopping material in the field of view, which lowers the proton flux at the detector but also increases the maximum proton energy to deflect. Those lessons are valid for any future mission requiring proton shielding. \\ From the present work, we identify soft proton scattering by X-ray focusing optics as one of the major sources of unwanted radiation in large effective area X-ray missions as the ATHENA X-ray telescope. The residual background level is proven to be orders of magnitude above the requirements, and a proton diverter is mandatory for the fulfillment of the mission science objectives. While these results come from a long activity of verification and physics validation of the scattering physics and Geant4 modeling, three main caveats are still present: (i) data from radiation environment observations of low energy protons start from $\sim50-60$ keV, and only extrapolations are possible at lower energies; (ii) laboratory measurements for proton scattering below 200 keV and at low $<1^{\circ}$ scattering angles are still missing, (iii) the proton diverter is assumed to be an ideal Halbach cylinder producing a uniform field within the magnet array. \\ In terms of accuracy in the input models, requiring a diverter energy threshold above 70 keV would efficiently deflect all protons below, despite their intensity. Solving once and for all the physics modeling for proton scattering at very long angle, very low energy is one of the aims of the ESA EXACRAD (Experimental Evaluation of Athena Charged Particle Background from Secondary Radiation and Scattering in Optics) project, where experimental measurements are currently ongoing. Geant4 simulations, using both Coulomb scattering as well as alternative new models (e.g. the Remizovich solution), will be compared to data and updates, if needed, will be provided to the community. Geant4 simulations of soft proton scattering by ATHENA mirrors will also be updated if necessary. Finally, in order to achieve a detailed and realistic simulation of the proton diverter, and comparing it with present assumptions, a synergy between the AREMBES and the SIMPOSIuM (SIMulations of Pore Optics in SIlicon and Modelling) ESA projects is currently ongoing, with the aim of using the output of the soft proton scattering simulation as input to a dedicated diverter simulator. | 18 | 8 | 1808.09431 |
1808 | 1808.10246_arXiv.txt | {Much of a planet's composition could be determined right at the onset of formation. Laboratory experiments can constrain these early steps. This includes static tensile strength measurements or collisions carried out under Earth's gravity and on various microgravity platforms. Among the variety of extrasolar planets which eventually form are (Exo)-Mercury, terrestrial planets with high density. If they form in inner protoplanetary disks, high temperature experiments are mandatory but they are still rare. Beyond the initial process of hit-and-stick collisions, some additional selective processing might be needed to explain Mercury. In analogy to icy worlds, such~planets might, e.g., form in environments which are enriched in iron. This requires methods to separate iron and silicate at early stages. Photophoresis might be one viable way. Mercury and {Mercury-like planets} % might also form due to the ferromagnetic properties of iron and mechanisms like magnetic aggregation in disk magnetic fields might become important. This review highlights some of the mechanisms with the potential to trigger Mercury formation.} \keyword{\textls[-5]{mercury; planet formation; protoplanetary disk; iron-silicate separation; extrasolar planets}} \begin{document} | Laboratory experiments are a valuable tool for testing and quantifying different mechanisms that might be important for planet formation. With a focus on Mercury, processes that act better at small distances to the star and that distinguish between iron and silicates are of special interest. Increasing temperatures are one aspect, and clearly, collisions depend on the temperature---its history as well as the actual temperature during collisions. No concluding remark can be given yet, except that growth changes. More work is needed in this field beyond the handful of experiments carried out. For the application of photophoresis, light, or at least temperature fluctuation, is a prerequisite. So, in a sense, this mechanism always has to fight for light. However, if any directed or fluctuating radiation field is present, at least ranging from visible to IR, its effect is huge, and it cannot be neglected. It has the advantage of combining two aspects: it is, in principle, stronger closer to the star and it can separate iron from silicates, both supporting the radial metal gradient and Mercury's iron nature. Last but not least, aggregation in magnetic fields also fulfills both requirements. Magnetic fields are stronger closer to the star, and aggregation is favoured for metallic iron dominated grains (below~1040 K). Laboratory experiments are limited to small size scales. So, these mechanisms can only trigger formation processes. How this is inherited by proceeding formation phases is a different story. {At~this point, it is unclear if this all regularly leads to the formation of Mercury planets or if Mercury is rather exceptional.} However, there are some promising candidate mechanisms supported by laboratory experiments to explain the high density of Mercury and the making of Exo-Mercury planets. \vspace{6pt} \funding{``This research received no external funding.'' } | 18 | 8 | 1808.10246 |
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1808 | 1808.05790_arXiv.txt | In the paradigm of magnetic acceleration of relativistic outflows, a crucial point is identifying a viable mechanism to convert the Poynting flux into the kinetic energy of the plasma, and eventually into the observed radiation. Since the plasma is hardly accelerated beyond equipartition, MHD instabilities are often invoked to explain the dissipation of the magnetic energy. Motivated by the fast variability that is shown by the gamma-ray flares of both AGN and PWNe, different authors have proposed the Poynting flux to be dissipated in a region where the flow is converging. Here we perform a linear stability analysis of ultra-relativistic, highly magnetised outflows with such a recollimation nozzle, showing that MHD instabilities are indeed induced by the convergence of the flow. The amplitude of the perturbations increases while recollimation gets stronger, and eventually diverges when the flow is focused to a single point. Hence, depending on the geometry of the outflow, instabilities excited while the flow is converging may play an important role to dissipate the magnetic energy of the plasma. | \label{sec:introduction} Relativistic jets occur ubiquitously among a wide variety of astrophysical sources, including Gamma Ray Bursts (GRBs), Active Galactic Nuclei (AGN), Pulsar Wind Nebulae (PWNe) and microquasars. It is widely accepted that these outflows are powered electromagnetically: the plasma slides on the magnetic field lines anchored to the rotating central object (a black hole or a neutron star) and is gradually accelerated by the magnetic tension (e.g. \citealt{Blandford1976, Lovelace1976, BlandfordZnajek1977}). In the aforementioned scenario the outflow's energy budget is initially dominated by the electromagnetic fields. Hence, it is crucial to understand how the Poynting flux is converted into the kinetic energy of the plasma, and eventually into the observed radiation. Since gradual acceleration is generally inefficient after equipartition between the magnetic and the kinetic energy is reached (e.g. \citealt{Komissarov2007, Komissarov2009, Lyubarski2009, Lyubarski2010, Lyubarski2011, Tchekhovskoy2008, Tchekhovskoy2009, Tchekhovskoy2010}), MHD instabilities are often invoked to destroy the regular structure of the flow and dissipate the magnetic energy (e.g. \citealt{Lyubarski1992, Eichler1993, Spruit1997, Begelman1998, Giannios2006}). The impact of the MHD instabilities depends on the causality condition in the lateral direction. In this paper we focus on the case of highly magnetised outflows, namely we assume the magnetisation to be $\sigma\gg 1$ (the Lorentz factor corresponding to the fast magnetosonic velocity is then $\sqrt{\sigma}$). If $\theta\gamma\gtrsim \sqrt{\sigma}$, the flow is causally disconnected in the lateral direction and global instabilities cannot develop \citep{Granot2011}. However, the plasma can pass through a weak shock whose only role is to make the downstream flow causally connected (for a more extended discussion see \citealt{Lyubarsky2012}). Hence, one can restrict the stability analysis to the causally connected regime $\theta\gamma\lesssim \sqrt{\sigma}$. If $\theta\gamma\lesssim 1$ (strong causal connection), the signal crossing time is shorter than the expansion time and the flow structure at any distance from the source is relaxed to an appropriate cylindrical equilibrium configuration \citep{Lyubarski2009}. The residual of the hoop stress and the electric force is balanced by the poloidal magnetic field and the ratio of the toroidal to poloidal components of the magnetic field is $B_\phi/B_{\rm p}\sim\gamma$, namely the two components are of the same order in the proper frame of the plasma. In the context of cylindrical jets, the most dangerous kink ($m=1$) modes have been investigated by a number of authors, both analytically (e.g. \citealt{IstominPariev1996, Lyubarski1999, Appl2000}) and numerically (e.g. \citealt{Nakamura2007, Mignone2010, Mizuno2012}). It turns out that the kink modes are stabilised if the strength of the poloidal magnetic field is nearly independent of the cylindrical radius $r$ (e.g. \citealt{IstominPariev1996, Lyubarski1999, Mizuno2012, Sobacchi2017}), which might be the case in a realistic scenario \citep{Narayan2009}. If $1\lesssim\theta\gamma\lesssim \sqrt{\sigma}$ (weak causal connection), the flow expands significantly within a signal crossing time and thus it is not in transverse equilibrium. However, the flow can still be adjusted to the external pressure without passing through a shock. One can show that the toroidal component of the magnetic field is much larger, in the proper frame, than the poloidal one \citep{Lyubarski2009}. Therefore one can consider the flow as composed of coaxial magnetic coils. It is important to realise that the confining pressure of the external medium can induce recollimation nozzles along the flow \citep{Lyubarski2009, GlobusLevinson2016}. Focusing on the case of PWNe, \citet{Lyubarsky2012} suggested how the relevant instability should work. Consider a region where the magnetic field can be effectively approximated as toroidal, so that the flow can be conceived as composed of coaxial magnetic coils, and take a small radial displacement of two coils. As the flow is converging the relative displacement increases, eventually becoming comparable with the radius of the coil itself. Hence, the regular structure of the field lines is destroyed and the recollimation nozzle provides a viable location to dissipate the magnetic energy. More recently, numerical simulations (\citealt{BrombergTchekhovskoy2016}; see also \citealt{Barniol2017}) have investigated a setup that might resemble the physical conditions of GRBs, namely a jet launched by a magnetar-like engine and collimated by the pressure of the stellar environment. Interestingly, it was found that (i) the jet generally experiences a rapid expansion that forces $B_\phi/B_{\rm p}\gg\gamma$, and is then recollimated by the external pressure; (ii) the instability indeed develops close to the nozzle. Due to the robustness of such a feature, these authors suggested the instabilities excited while the flow is converging to be the mechanism that powers the GRB prompt emission. The case of dissipation occurring in the vicinity of a recollimation nozzle is supported by the fact that the gamma-ray emission is often variable on extremely short time scales ($\lesssim\text{day}$ for both typical AGN and PWNe flares, and even down to few minutes in the former case), which poses serious constraints on the size of the emitting region. For this reason, different authors have proposed that the observed flares are produced in a region where the flow is converging, so that its transverse size is reduced; see for example \citet{BrombergLevinson2009, GlobusLevinson2016} in the context of AGN and \citet{Lyubarsky2012} in the context of PWNe, and references therein. Here we present a linear stability analysis of converging outflows far from a local cylindrical equilibrium, aiming to identify the fundamental physical parameters that regulate the growth of the instability. We find that the growth of the unstable modes is simply determined by the geometry of the outflow, with narrower recollimation nozzles being more unstable, while it is independent of other features such as the Lorentz factor of the flow. The paper is organised as follows. In Section \ref{sec:equations} we present the fundamental equations governing Poynting-dominated outflows and we find a stationary solution for their structure. In Section \ref{sec:evolution} we study the linear evolution of perturbations propagating along the flow. Finally, in Section \ref{sec:conclusions} we summarise our conclusions. | \label{sec:conclusions} We have performed a linear stability analysis of ultra-relativistic, highly magnetised outflows with a recollimation nozzle. We have considered perturbations propagating through the region close to the nozzle, where the flow is approximately conical. We have neglected the dynamical effect of the poloidal magnetic field, as appropriate in the causally disconnected regime. We have furthermore worked in the limit $kR\ll\gamma$, where $k$ is the wavenumber of the perturbation, $R$ is the outer radius and $\gamma$ is the Lorentz factor of the flow. This corresponds to the case where the wavelength of the perturbation is longer than $R$ in the proper frame of the flow. Our main results can be summarised as follows: \begin{enumerate} \item while $R$ shrinks from $R_1$ to $R_2<R_1$, the amplitude of non-axisymmetric perturbations propagating along the flow grows by a factor $R_1/R_2$ with respect to the unperturbed fields. This confirms the simple picture proposed by \citet{Lyubarsky2012}. Consider a purely toroidal magnetic field, so that the flow can be conceived as composed of coaxial magnetic coils. If one takes a small radial displacement $\delta R$ of two coils, the relative amplitude of the perturbed magnetic field is $\delta B/B \sim \delta R/R$. Assuming that $\delta R$ remains approximately constant as the outer radius $R$ shrinks, $\delta B/B$ indeed grows proportionally to $R^{-1}$. \item the growth of the perturbations depends only on the geometry of the flow (namely $R_1$ and $R_2$), being instead independent of its Lorentz factor. Since the perturbations are amplified by an arbitrarily large factor in the limit $R_2\ll R_1$, the flow becomes violently unstable if the recollimation is strong enough. Instabilities excited while the flow is converging may therefore play an important role to dissipate the magnetic energy of the plasma, and understanding if/how much a realistic outflow is recollimated is crucial to assess their impact. \end{enumerate} In order to achieve a strong recollimation, the outflow needs first to get rid of the poloidal field that halts the compression through its magnetic pressure. The dynamical effect of the poloidal field is determined by its ratio with the toroidal field in the proper frame of the flow, namely $B_{\rm p}$ is negligible if $B_\phi/\gamma B_{\rm p}\gg 1$. Since $B_\phi/\gamma B_{\rm p}$ is proportional to $R^2$, a fast initial expansion allows the outer radius to shrink significantly thereafter. The amount of initial expansion is clearly connected to the properties of the central engine and to those of the confining medium. In general, we expect a significant expansion to be favoured when the initial (mostly magnetic) pressure of the ejected plasma largely exceeds that of the confining medium. In the case of PWNe, where the pressure of the confining medium is extremely low, the instability studied in this paper is therefore a promising explanation for the observed gamma-ray flares. We refer to \citet{Lyubarsky2012} for a more extended discussion of this point in the context of the Crab nebula. In the case of GRBs, the jet may expand significantly before being recollimated by the pressure of the stellar environment if it is launched by a magnetar-like engine (see for example the simulations of \citealt{BrombergTchekhovskoy2016}). Hence, as suggested by these authors, also in these objects the instabilities excited while the flow is converging may be responsible for the dissipation of the magnetic energy. This claim is supported by the fact that we find these instabilities not to be suppressed even in the ultra-relativistic regime. Finally, note that we have not addressed the question of how the magnetic energy is actually converted into the observed radiation. Indeed, the dissipation of the magnetic energy is an intrinsically non-linear process whose study would require numerical simulations, which is out of the scope of the paper. Instead, we have simply assumed that the onset of MHD instabilities triggers the dissipation of the magnetic energy destroying the regular structure of the field lines. The actual dissipation may happen via reconnection and/or dissipation of MHD turbulence, as suggested by the results of numerical simulations (see for example \citealt{BrombergTchekhovskoy2016}). | 18 | 8 | 1808.05790 |
1808 | 1808.05045_arXiv.txt | In this paper the focus is on inflationary dynamics in the context of Einstein Gauss-Bonnet gravitational theories. We investigate the implications of the slow-roll condition on the slow-roll indices and we investigate how the inflationary dynamical evolution is affected by the presence of the Gauss-Bonnet coupling to the scalar field. For exemplification of our analysis we investigate how the dynamics of inflationary cubic, quartic order and also exponential scalar potentials are affected by the non-trivial Gauss-Bonnet coupling to the scalar field. As we demonstrate it is possible to obtain a viable phenomenology compatible with the observational data, although the canonical scalar field theory with cubic and quartic order potentials does not yield phenomenologically acceptable results. In addition, with regard to the exponential potential example, the Einstein Gauss-Bonnet extension of the single canonical scalar field model has an inherent mechanism that can trigger the graceful exit from inflation. Furthermore we introduce a bottom-up reconstruction technique, in the context of which by fixing the tensor-to-scalar ratio and the Hubble rate as a function of the $e$-foldings number, one is capable of reproducing the Einstein Gauss-Bonnet theory which generates the aforementioned quantities. We illustrate how the method works by using some relatively simple examples. | One of the mysteries of the primordial evolution of our Universe is related to the inflationary era. The theoretical mechanism called inflation was initially introduced in the early 80's \cite{Guth:1980zm,Starobinsky:1982ee,Linde:1983gd} in order to solve the homogeneity, isotropy, flatness and horizon problems that haunted the standard Big Bang model of cosmology. The standard inflationary paradigm is realized by using a canonical scalar field dubbed inflaton, the quantum fluctuations of which may explain the Cosmic Microwave Background anisotropies. With regard to the latter, the latest Planck observations \cite{Ade:2015lrj} have constrained the power spectrum of the primordial curvature perturbations to be nearly scale invariant. Up to date, quite many inflationary models have been introduced in the literature, some of which remain viable to a great extent even after the Planck observations \cite{Ade:2015lrj}, for reviews see \cite{inflation1,inflation2,inflation4,reviews1,reviews2}. We should note that the inflationary paradigm is not the only mechanism which may generate a nearly scale invariant power spectrum, since the latter can be achieved even in the context of bouncing cosmology \cite{bounce}. In this paper we shall confine ourselves on the study of slow-roll inflationary models which are string inspired, namely an Einstein Gauss-Bonnet inflationary models. Many studies of this sort appear in the literature, and for an important stream of related papers see for example \cite{Nojiri:2006je,Calcagni:2005im,Calcagni:2006ye,Cognola:2006sp,Nojiri:2005vv,Nojiri:2005jg,Nojiri:2007te,Bamba:2014zoa,Yi:2018gse,Guo:2009uk,Guo:2010jr,Jiang:2013gza,Koh:2014bka,Koh:2016abf,Kanti:2015pda,vandeBruck:2017voa,Kanti:1998jd,Nozari:2017rta,Chakraborty:2018scm}. The study of the Einstein Gauss-Bonnet inflationary models is highly motivated from the fact that the inflationary era occurs very close to the primordial era of our Universe which is believed to be described by quantum gravity. The complete description of this primordial quantum era is not known up to date, but various aspects of string theory can provide useful insights regarding this era. Then it is highly possible that a residual imprint of the quantum Universe onto the classical inflationary Universe can be found and it may have an observable effect on the primordial power spectrum. Therefore, the low energy effective theory that results from the primordial quantum gravity theory may affect the inflationary era directly. In this respect, the Einstein Gauss-Bonnet theory is an appealing candidate theory for the inflationary era, since it is obtained by a string theory. Particularly, the underlying string theory induces a curvature correction to the classical canonical scalar field theory, which depends on the Gauss-Bonnet invariant $\mathcal{G}$. Thus, among the various modified gravities \cite{reviews1,reviews2,reviews3,reviews4,reviews5,reviews6}, the Einstein Gauss-Bonnet theory is quite interesting since it originates from string theory. In view of these appealing features of Einstein Gauss-Bonnet theories, in this paper we shall study the inflationary era produced by these theories in the slow-roll approximation. Firstly, we shall investigate how the slow-roll condition is realized in the context of these theories, and we shall calculate explicitly the slow-roll indices of the theory at hand. We shall also study several classes of inflationary models and we will calculate the observational indices of inflation, which we shall compare with the observational data coming from Planck. As we will demonstrate, the resulting theory can be compatible with the observations for a wide range of parameters. Also the second part of this paper is devoted on studying reconstructed slow-roll models of Einstein Gauss-Bonnet gravity from the observational indices. This bottom-up approach is frequently used in the literature \cite{Odintsov:2018ggm,Odintsov:2017fnc,Narain:2017mtu,Narain:2009fy,Fumagalli:2016sof,Gao:2017uja,Lin:2015fqa,Miranda:2017juz,Fei:2017fub,Herrera:2018mvo,Herrera:2018ker} and it provides viable inflationary models. We shall provide analytic formulas for the slow-roll indices and for the spectral index of the primordial curvature perturbations and we will investigate the viability of the resulting models in detail. An similar reconstruction scheme to the one we shall develop, was presented by A. Starobinsky in \cite{starobconf}. We should note that in the article we shall use the general formalism developed in Refs. \cite{Noh:2001ia,Hwang:2005hb,Hwang:2002fp,Kaiser:2013sna}, however we shall extend it in the slow-roll approximation, and moreover we shall investigate the behavior of specific models. The new features of this work are the development of the slow-roll inflation formalism in the context of Einstein Gauss Bonnet theories and the introduction of the bottom-up reconstruction techniques from the observational indices, developed for the aforementioned theories. This paper is organized as follows: In section II we present the essential features of Einstein Gauss-Bonnet theory and we investigate how the slow-roll condition modifies the dynamics of inflation, in terms of the slow-roll indices. We present analytic expressions for the slow-roll indices in the slow-roll approximation and we calculate the observational indices of inflation, namely the spectral index of the primordial curvature perturbations and the tensor-to-scalar ratio. In addition, we use two illustrative examples in order to exemplify our results. In section III we introduce a bottom-up reconstruction technique in which given the tensor-to-scalar ratio and the Hubble rate, one is able to find which theory realizes the aforementioned quantities. We provide analytic formulas for the slow-roll indices and for the spectral index, and we exemplify our technique by using a characteristic example. Finally the conclusions follow in the end of the paper. Prior starting, it is worth to describe in brief the geometric background assumed in this paper, which is a flat Friedmann-Robertson-Walker (FRW) background, with line element, \begin{equation} \label{metricfrw} ds^2 = - dt^2 + a(t)^2 \sum_{i=1,2,3} \left(dx^i\right)^2\, , \end{equation} where $a(t)$ is the scale factor. In addition, the metric connection is the Levi-Civita, which is torsion-less, symmetric and metric compatible. | In this paper the focus was on inflationary solutions of Einstein Gauss-Bonnet gravity. Particularly we investigated the implications of the slow-roll conditions on the gravitational theory by imposing the condition that the slow-roll indices are much smaller than unity during the slow-roll era. In turn the slow-roll condition restricted the gravitational theory quantified by the coupling of the Gauss-Bonnet scalar $\xi (\phi)$ and by the scalar potential $V(\phi)$. After discussing the general formalism, we investigated the phenomenological implications of power-law and exponential potentials focusing on the cubic, quadratic order potentials and simple exponential potentials. The power-law potentials in the context of a single canonical scalar field theory fail to comply with the Planck 2015 observational data, however as we demonstrated, in the context of Einstein Gauss-Bonnet theory it is possible to obtain a phenomenologically viable theory, and in fact for a wide range of values of the free parameters of the theory. Hence a first outcome of this work is that the Einstein Gauss-Bonnet theories can make non-viable single canonical scalar field theories to be viable, and actually the coupling of the scalar field to the Gauss-Bonnet scalar controls this procedure. With regard to the exponential potentials, the classical single scalar theory has no inherent mechanism to trigger the graceful exit from inflation, since the slow-roll indices are constant and field-independent. However the Einstein Gauss-Bonnet version of the theory has the slow-roll index $\epsilon_4$ which is field dependent, and thus the slow-roll phase ends when this index becomes of order $\mathcal{O}(1)$. In the second part of this work we presented a bottom-up reconstruction method in the context of which it is possible to obtain a viable inflationary theory by fixing the functional form of the observational indices, and particularly of the tensor-to-scalar ratio. Specifically, one fixes the Hubble rate and the tensor-to-scalar ratio as functions of the $e$-foldings number. After that, by appropriately expressing the slow-roll indices and the equations of motion as functions of the $e$-foldings number, it is possible to obtain the exact Einstein Gauss-Bonnet theory that may realize the given Hubble rate and the tensor-to-scalar ratio, provided that the slow-roll conditions hold true. Particularly, one obtains the functions $\xi (\phi (N)))$, and $\phi (N)$, and from these the potential $V(N)$ is obtained. By using the resulting functions, the calculation of the spectral index of the primordial curvature perturbations is straightforward, and hence it can be checked when the resulting theory is compatible with the observational data. Finally, if the resulting function $\phi (N)$ can be inverted, the scalar potential $V(\phi)$ and the function $\xi (\phi)$ can be obtained. In principle, the methodology we introduced in this work can be applied in generalized forms of Einstein Gauss-Bonnet gravity, in which case the gravitational part containing the scalar curvature $R$ may be a function of $R$, that is, an $f(R)$ gravity, or it may be a function of both the scalar curvature and of the scalar field of the form $f(R,\phi)$. In this case the resulting equations of motion might be more complicated, however the slow-roll conditions on the indices may simplify the calculations to a great extent. We hope to address soon some of the above issues in a future work. | 18 | 8 | 1808.05045 |
1808 | 1808.02981_arXiv.txt | The magnesium-potassium anti-correlation observed in globular cluster NGC2419 can be explained by nuclear burning of hydrogen in hot environments. The exact site of this nuclear burning is, as yet, unknown. In order to constrain the sites responsible for this anti-correlation, the nuclear reactions involved must be well understood. The \Kpp reactions are one such pair of reactions. Here, we report a new evaluation of the \pg reaction rate by taking into account ambiguities and measurement uncertainties in the nuclear data. The uncertainty in the \pg reaction rate is larger than previously assumed, and its influence on nucleosynthesis models is demonstrated. We find the \pg reaction cross section should be the focus of future experimental study to help constrain models aimed at explaining the magnesium-potassium anti-correlation in globular clusters. | \label{sec:intro} The alkali element potassium is synthesized in several stellar environments. It is predominately produced in a combination of hydrostatic and explosive oxygen burning~\cite{WW_2002}, conditions found only in highly-evolved massive stars during the pre-supernova phase and the ensuing explosion. However, models of galactic chemical evolution, which are based on the nucleosynthetic yields of supernovae, so far severely under-predict the observed potassium abundance in our galaxy~\cite{Timmes_1995,Goswami_2000,Romano_2010,Prantzos2018}. Potassium is also synthesized in smaller quantities in high-temperature hydrogen burning environments, which are believed to be important in explaining elemental abundance signatures in globular clusters. Of particular interest is NGC 2419~\cite{Mucciarelli_2012}, where it was found that a significant fraction of its member stars ($\approx40\%$) are highly enriched in elemental potassium. Additionally, there is a strong anticorrelation observed between potassium and magnesium abundances, reminiscent of the ubiquitous Na-O and Mg-Al anticorrelations found much more commonly in clusters (see Ref. \cite{Gratton_2012} and references therein). Though the main reason for these discrepancies has not been established, a more accurate description of potassium synthesis will be helpful in this area. To achieve this, the potassium destruction reactions \Kpp are crucial. Iliadis \textit{et al.} \cite{Iliadis2016} explored the astrophysical conditions that could be responsible for isotopic correlations in NGC 2419. Their method featured a Monte Carlo nucleosynthesis network that included all known uncertainties in the thermonuclear reaction rates. They obtained a range of stellar temperatures and densities that quantitatively reproduced all of the elemental abundances measured in the potassium-rich stars. Later, Ref. \cite{Dermigny_2017} extended that method by including a sensitivity study of the nuclear reaction rates. They found several reactions whose rates need to be better constrained in order to more accurately identify an astrophysical site responsible for the anomalies. The majority of these pertain to the synthesis and destruction of \nuc{39}{K}. The rate of one of these reactions, \pg, was based on preliminary calculations, so it is important to reinvestigate the $^{39}$K $+$ p reactions based on a full evaluation of the nuclear physics input. In this paper, we calculate the rate of the \pg reaction using all available experimental information. The \pa rate was found to not significantly influence final abundances in the stellar environments of interest here, so we leave its evaluation to future work. In Sec. \ref{sec:formalism}, a brief overview of the reaction rate formalism is presented along with the Monte Carlo method used to calculate uncertainties on the rate given experimental uncertainties on the cross sections. In Sec.~\ref{sec:rates}, details of the experimental information are presented. The Monte Carlo rates using that information are computed in Sec.~\ref{sec:rates-results} and compared to previous reaction rate calculations. Astrophysical implications of these rates as they pertain to nucleosynthesis in globular clusters are presented in Sec.~\ref{sec:astro-results}, and all is summarized in Sec.~\ref{sec:conclusions}. | \label{sec:conclusions} The \pg reaction has been found previously to affect potassium synthesis in stellar environments leading to the Mg-K anticorrelation in the globular cluster NGC 2419~\cite{Dermigny_2017}. That finding was based on estimates of the current experimental uncertainty of the reaction cross sections, which spurred a thorough re-investigation of the current experimental picture. By considering current experimental measurements of narrow resonances and including full characterization of upper limits on unobserved resonance strengths, we present here updated estimates of the rate of the \pg reaction. The former reaction rate uncertainties also include ambiguities between experimentally determined resonance strengths reported in Refs.~\cite{Leenhouts1966}, \cite{Cheng1981}, and~\cite{Kikstra1990}. Correlations between measurements are also taken into account. The results of this investigation show that the uncertainties in the \pg reaction are larger than previously estimated. The nucleosynthesis ramifications of these findings are also presented by considering an astrophysical scenario within the bounds established in Ref.~\cite{Dermigny_2017}. We find that the increased uncertainty in the \pg reaction rate establishes a clear correlation between it and the final abundance of \nuc{39}{K}. Furthermore, the predicted uncertainty in the elemental abundance of potassium is broadened towards lower values. Clearly, the \pg reaction must be better measured if astrophysical scenarios explaining the Mg-K anti-correlation are to be constrained. | 18 | 8 | 1808.02981 |
1808 | 1808.03455_arXiv.txt | The spatial distribution of known globular clusters (GCs) in the Milky Way shows that the current census of GCs is incomplete in the direction of the Galactic plane. We present the discovery of two new GCs located close to the Galactic plane in the sky. These two GCs, RLGC 1 and RLGC 2, were discovered serendipitously during our new cluster survey \citep{ryu18} based on near-Infrared and mid-Infrared survey data. The two GCs show a grouping of resolved stars in their $K$ band images and a presence of faint diffuse light in their outer regions in the \textit{WISE W1} band images. They also show prominent red giant branches (RGBs) in their $K$ \replaced{--}{vs.} $(J-K)$ color-magnitude diagrams (CMDs). We determine structural parameters of the two GCs using King profile fitting on their $K$ band radial number density profiles. The determined values are consistent with those of known GCs. Finally, we determine the distances, metallicities, and reddenings of the two GCs using the isochrone fitting on their CMDs. For the fitting, we assume that the ages of the two GCs are 12.6 Gyr and the brightest RGB stars of each cluster correspond to the tip of the RGB. Distances and metallicities of the two GCs are estimated to be $d=28.8\pm4.3$ kpc and $\textrm{[Fe/H]}=-2.2\pm0.2$ for RLGC 1 and $d=15.8\pm2.4$ kpc and $\textrm{[Fe/H]}=-2.1\pm0.3$ for RLGC 2. These results show that the two GCs are located at the far-half region of the Milky Way and they may belong to the halo of the Milky Way. | The current census of the Milky Way Globular Clusters (GCs) is incomplete in the direction of the Galactic plane. The severe extinction of the Galactic plane prevents us from finding GCs located in the far-half region of the Milky Way (behind the center of the Milky Way). % According to the catalog of the Milky Way GCs in \citet[2010 edition]{har96}, there are 57 GCs at $|Z|<1$ kpc. While 43 of these GCs are in the close-half of the Galaxy, only 14 GCs are in the far-half region. In addition to the GCs in the Harris catalog, many new GCs \deleted{and GC candidates} in the Milky Way were discovered since 2005. The total number of these objects amounts to \replaced{117}{55} :Willman 1 \citep{wil05}, FSR 584 \citep{bic07}, FSR 1767 \citep{bon07}, FSR 190 \citep{fro08a}, FSR 1716 \replaced{\citep{fro08b}}{\citep[=VVV--CL005;][]{fro08b, min17a}}, Pfleiderer 2 \citep{ort09}, SEGUE 3 \citep{bel10}, Mercer 5 \citep{lon11}, VVV--CL001 \citep{min11}, VVV--CL002 and CL003 \citep{mon11}, Mu\~noz 1 \citep{mun12}, Kronberger 49 \citep{ort12}, Balbinot 1 \citep{bal13}, VVV--CL110, CL128, CL131, CL143, and CL150 \citep{bor14}, Crater \citep{lae14, wei16}, Eridanus III \citep{bec15}, Kim 1 \citep{kim15a}, Kim 2 \citep{kim15b}, Laevens 3 \citep{lae15}, DES 1 \citep{luq16}, Kim 3 \citep{kim16}, Gaia 2 \citep{kop17}, Minniti 01--22 \citep{min17b}, Sagittarius II \citep{lae15, mut18}, \replaced{and Camargo 1102--1106 \citep{cam18}}{Camargo 1102 \citep{bic18, cam18}, and Camargo 1103--1106 \citep{cam18}}. Only $\sim30$ % among these new objects (VVV--CL clusters, Minniti 01--22, and Camargo 1102--1106) are confirmed to be located in % the central Galactic plane region. % Even if we include these objects, the numbers of the GCs in the close-half and far-half regions at $|Z|<1$ kpc are \replaced{48 and 30}{58 and 19}, respectively. This implies that there are more undiscovered GCs in the far-half region of the Galaxy. Recently, \citet{ryu18} carried out a new survey of star clusters in the Galactic central region ($|l|<30\arcdeg$ and $|b|<6\arcdeg$) using near-Infrared (NIR) surveys and mid-Infrared (MIR) surveys, such as the Two Micron All Sky Survey (2MASS; \citealt{skr06}) and the \textit{Wide-field Infrared Survey Explorer}(\textit{WISE}; \citealt{wri10}). They found 923 new star clusters. During this survey, we serendipitously discovered two new GC candidates at $(l, b)=(336\arcdeg.87, 4\arcdeg.30)$ and $(27\arcdeg.63, -1\arcdeg.04)$: Ryu 059 and Ryu 879 (called RLGC 1 and RLGC 2 hereafter), which are reported in this Letter. This paper is organized as follows. We introduce the selection criteria for the GC candidates in Section 2. The two clusters turn out to be old GCs, thus we derive their structural parameters and distances, metallicities, and reddenings in Section 3. In Section 4, we discuss and compare spatial locations and physical parameters of the new GCs with those of known GCs. Finally, we conclude with a number estimation of undiscovered GCs in the far-half region of the Galactic plane. | \subsection{Spatial Locations} \begin{figure*} \epsscale{0.8} \plotone{fig3.eps} \caption{(a) Spatial distribution of RLGC 1, RLGC 2, and the known GCs on the face-on view of the Milky Way Galaxy \citep{chu09}. Dashed lines are guide lines for the location of the Galactic center. The Sun is represented as the Sun symbol. Metallicities of GCs are color-coded as shown in the color bar at the top. The sizes of the symbol represent relative magnitudes of GCs. RLGC 1 and RLGC 2 are emphasized by the star symbol, irrespective of their magnitudes. Solid lines represent distance errors of the clusters. Open triangles represent recently reported GCs (\citealt{min17b} and \citealt{cam18}). (b) Spatial distribution of GCs on the edge-on view of the Milky Way. Symbols are the same as those in (a). The $|Z|<1$ kpc region is represented as the yellow shaded region. \label{spd}} \end{figure*} In Figure \ref{spd} we plot the spatial location of the new GCs in comparison with other known GCs listed in \citet{har96}, \replaced{\citet{cam18}, and Piatti (2018)}{\citet{min17b}, and \citet{cam18}}, in both the face-on view (Figure \ref{spd}(a)) and the edge-on view (Figure \ref{spd}(b)). \added{We adopted the distances to Minniti 01--22 given by \citet{min17b} (see \citet{pia18} for other distance estimation)}. The two GCs are located at the far-half region of the Milky Way. Their distances from the closest neighbor GCs are 10.2 kpc (NGC 5824) and 4.5 kpc (Pal 11) for RLGC 1 and RLGC 2; practically, no neighbor GCs are found in the vicinity the new GCs. Based on the distances and Galactic latitudes of the new GCs, we derive their vertical positions from the Galactic plane: $Z=2.2\pm0.3$ kpc for RLGC 1 and $Z=-290\pm40$ pc for RLGC 2. % RLGC 1 that has a low metallicity is likely to be a halo GC, located above the thick disk \citep[$h_z=0.9\pm0.1$ kpc]{li17}. RLGC 2 is located in the thick disk. However, its low metallicity ([Fe/H]$=-2.1\pm0.3$) indicates that it must be a halo GC. Therefore, RLGC 2 may be a halo GC passing through the disk now. \subsection{Absolute Magnitudes} \citet{coh07} provided the integrated $K_s$ magnitude of the known GCs derived with $50\arcsec$ radius apertures. For comparison, we derive $50\arcsec$-integrated $K$ magnitudes of the two GCs. The $50\arcsec$-integrated magnitudes are $K_{50\arcsec}=-10.28\pm0.30$ mag (RLGC 1) and $K_{50\arcsec}=-10.10\pm0.31$ mag (RLGC 2), while the peak absolute magnitude of the GC luminosity function noted in \citet{coh07} is $M_K =-9.7$ mag. The magnitudes of the new GCs are 0.4--0.6 mag brighter than the peak magnitude of the known GCs in the Milky Way. Using the relation for the metal-poor GCs in \citet{coh07}: $V-K_s=2.93+0.409$[Fe/H], we estimate the $50\arcsec$-integrated $V$ magnitudes of the new GCs. The estimated magnitudes are $V_{50\arcsec}=-8.25\pm0.31$ mag (RLGC 1) and $V_{50\arcsec}=-8.03\pm0.33$ mag (RLGC 2), while the peak absolute magnitude % is $M_V=-7.66\pm0.09$ mag \citep{dic06}. \deleted{The integrated $V$ magnitudes are also 0.4--0.6 mag brighter than the peak magnitude of the known GCs.} \subsection{Structural Parameters} The core radii of the two GCs are $r_c=1.51\pm0.34$ pc and $r_c=0.97\pm0.15$ pc for RLGC 1 and RLGC 2. These values are consistent with the median core radius of the known GCs, which is $median(r_c)=1.04$ pc. The half-light radii of the two GCs are also comparable with the known GCs: $r_h=4.59\pm0.74$ pc for RLGC 1, $r_h=2.14\pm0.33$ pc for RLGC 2, and $median(r_h)=3.03$ pc. The concentration indices of the new GCs are relatively lower than those of the known GCs: $c=0.7\pm0.2$ for RLGC 1 and $c=0.6\pm0.1$ for RLGC 2 ($cf.\; median(c)=1.50$). The core radii of the new GCs are consistent with those of known GCs, hence low concentration indices would be related to tidal radii. Actually, the tidal radii of the two GCs ($r_t=7.79\pm3.55$ pc for RLGC 1 and $r_t=4.14\pm0.78$ pc for RLGC 2) are much smaller than the median tidal radius of the known GCs, $median(r_t)=28.86$ pc. This implies that the derived tidal radii of the new GCs might have been underestimated. However, complex backgrounds in the data we used prevent us from recognizing weak enhancements of number densities in outer cluster regions. \subsection{Concluding Remarks} Based on their morphologies, radial number density profiles, CMDs, and other determined parameters, RLGC 1 and RLGC 2 are likely to be genuine metal-poor halo GCs. However, our photometric parameters are based on uncertain assumptions: The magnitude of the brightest RGB star and the selective extinction value $R_V$. These assumptions are likely to be the origin of additional distance uncertainties. Deeper NIR photometry reaching the main-sequence turn-off point is needed to measure more accurate distances, ages, and metallicities of the two new GCs. \added{ Considering the low metallicity of these clusters, we expect that blue horizontal branch stars might be detected in deep optical CMDs reaching fainter than $V\sim23$ mag.} Including these two new GCs, the current census of the GCs \deleted{and candidates} in the Milky Way is \replaced{$N_{\textrm{GC}}=276$}{$N_{\textrm{GC}}=214$}. The current GC numbers at $|Z|<1$ kpc are \replaced{48 and 31}{58 and 20} in the close-half region and the far-half region, respectively. Therefore, we expect that there are about \replaced{20}{30} % undiscovered GCs in the far-half and $|Z|<1$ kpc region of the Galactic disk. | 18 | 8 | 1808.03455 |
1808 | 1808.04934_arXiv.txt | The solar eruption on 2012 January 27 resulted in a wide-spread solar energetic particle (SEP) event observed by \stereoa{} and the near-Earth spacecraft (separated by 108\degree). The event was accompanied by an X-class flare, extreme-ultraviolet (EUV) wave and fast coronal mass ejection (CME). We investigate the particle release by comparing the release times of particles at the spacecraft and the times when magnetic connectivity between the source and the spacecraft was established. The EUV wave propagating to the magnetic footpoint of the spacecraft in the lower corona and the shock expanding to the open field line connecting the spacecraft in the upper corona are thought to be responsible for the particle release. We track the evolution of the EUV wave and model the propagation of the shock using EUV and white-light observations. No obvious evidence indicates that the EUV wave reached the magnetic footpoint of either \stereoa{} or L1-observers. Our shock modeling shows that the release time of the particles observed at L1 was consistent with the time when the shock first established contact with the magnetic field line connecting L1-observers. The release of the particles observed by \stereoa{} was delayed relative to the time when the shock was initially connected to \stereoa{} via the magnetic field line. We suggest that the particle acceleration efficiency of the portion of the shock connected to the spacecraft determines the release of energetic particles at the spacecraft. | A straightforward interpretation for a gradual solar energetic particle (SEP) event observed simultaneously by spacecraft at different locations in the heliosphere is the wide extent of a coronal mass ejection (CME)-driven shock in the corona or interplanetary (IP) medium \citep[e.g.,][]{Cliver95, Heras95, Reames99, Reames13}. The particles are accelerated at the shock and then injected onto the magnetic field lines connecting the observers in the heliosphere \citep[e.g.,][]{Cliver04, Zank07,Kozarev15}. Cross-field diffusion processes have also been suggested to account for the particle transport over a wide longitudinal range from a narrow particle source region \citep[e.g.,][]{Wang12,Lario17_Aug14}. Studies based on the joint observations from the \emph{Solar-TErrestrial RElations Observatory} \citep[\stereo;][] {Kaiser08} and near-Earth spacecraft such as the \emph{Advanced Composition Explorer} \citep[\ace;][] {Stone98} and the \emph{Solar and Heliospheric Observatory} \citep[\soho;][] {Domingo95}, have been carried out for the longitudinal distribution and the release mechanism of energetic particles \citep[e.g.,][]{Rouillard12,Prise14, Richardson14, Gomez-Herrero15,Kouloumvakos15,Kozarev16,Rouillard16,Zhu16, Lario13, Lario14, Lario16, Lario17_Aug14, Kwon17}. In addition, the remote-sensing observations from the \emph{Atmospheric Imaging Assembly} \citep[AIA;][] {Lemen12} on board the \emph{Solar Dynamics Observatory} \citep[\sdo;][] {Pesnell12} with high cadence and from the \emph{Sun-Earth Connection Coronal and Heliospheric Investigation} \citep[SECCHI;][] {Howard08} instruments on board \stereo{} offer an unprecedented opportunity to investigate the EUV wave and shock evolution and the connectivity between the observers and the coronal disturbances. The properties of the CME-driven shock and accompanied physical phenomena are used to study the mechanisms of particle acceleration and release and hence explain the wide distribution of SEP events. EUV waves, which propagate in the solar corona and can be detected in EUV images, are considered as one of the candidates to be responsible for the particle release in the SEP events. \citet{Rouillard12} examine the SEP event on 2011 March 21 and conclude that the arrival times of the EUV wave at the magnetic footpoints of the spacecraft coincide with the particle release times. \citet{Park13} suggest that the EUV wave traces the release site of SEPs accelerated by the CME-driven shock in a study of 12 SEP events. Statistical work by \citet{Miteva14} with 179 SEP events show that a large majority of SEP events are associated with EUV waves. They get a connectivity between the extrapolated arrival times of the EUV wave at the magnetic footpoint of Earth and the particle onsets for 26 eastern SEP events. The other candidate to explain the release of particles is the expansion of the shock in the outer corona. It is suggested that the particle release times are associated with the times when the shock establishes contact with the magnetic field lines connecting the spacecraft, rather than the arrival times of the EUV wave at the magnetic footpoints of the spacecraft. \citet{Prise14} and \citet{Gomez-Herrero15} study the SEP event on 2011 November 3 with different methods. They conclude that the EUV wave is too slow to explain the particle release, the expansion of the CME-driven shock at higher altitudes is consistent with the release times of particles at different spacecraft. \citet{Kwon14} develop a geometrical model to determine the three-dimensional (3D) structure of the shock using the EUV and white-light (WL) observations from multipoint spacecraft. \citet{Lario16} apply the geometrical model to fit the structure of the shock and indicate that the particles are released when the portion of the shock magnetically connected to each spacecraft is at a relatively high altitude. Other physical mechanisms may also play a role in the particle release and the wide longitudinal extent of an SEP event. \citet{Kouloumvakos16} discuss the role of the cross-field diffusion for the particle transport by tracking the EUV wave and modeling the shock on 2012 March 7. In this paper, we study the release of energetic particles using coordinated imaging and in situ observations from \stereo{} and the near-Earth spacecraft. The roles of the associated EUV wave and shock during the release of particles are still under debate. It is necessary to investigate this problem using combined imaging and in situ data. We focus on the 2012 January 27 solar eruption that originated from NOAA AR 11402 (N29\degree W86\degree). It was associated with an X1.7 flare starting at 17:37 UT and peaking at 18:37 UT. It produced an intense SEP event accompanied by a wide expanded EUV wave, halo CME and CME-driven shock. The CME speed was about 2508 km s${{}^{-1}}$ estimated in the \soho/LASCO CME catalog (\url{https://cdaw.gsfc.nasa.gov/CME_list/}). The center panel of Figure 1 shows the longitude of the active region (red arrow) and the positions of \stereoa, ${B}$ and Earth during this event. \stereoa{} was located at 22\degree{} west of the active region, \stereob{} was located at the opposite side of the Sun, and Earth was 86\degree{} east with respect to the active region. These widely separated spacecraft offer EUV and WL observations for the associated EUV wave, CME and shock from three vantage points. The active region was imaged by \stereoa{} on the solar disk, while the near-Earth spacecraft were in a good position to offer lateral imaging observations and measure the energetic particles (see Figure 1). This configuration provides a good opportunity to track the evolution of the EUV wave in the corona and the shock at higher altitudes as well as to determine their connection with the particle release. In Section 2, we describe the observations and perform the analysis. The results are summarized and discussed in Section 3. | The solar eruption on 2012 January 27 was associated with an X1.7 flare, EUV wave, CME-driven shock and wide-spread SEP event. Combining multipoint remote sensing and in situ observations, we have estimated the release times of the particles observed by the spacecraft and the locations of the magnetic footpoints of the spacecraft, tracked the evolution of the EUV wave along the directions from the active region to the magnetic footpoints of the spacecraft, modeled the shock and identified the magnetic connectivities between the shock and the spacecraft at the estimated particle release times. The results are summarized and discussed as follows. Near-Earth spacecraft and \stereoa{} observed the shock in situ and the SEP event associated with the January 27 solar event, which is a backsided event for \stereob{}. The high-energy particle intensities at both L1-observers and \stereoa{} exhibited almost immediate onsets, obvious increases and long durations above the pre-event intensity levels. The particle intensities at \goes{} peaked quickly after the onset and those at \stereoa{} peaked around the shock passage with two orders of magnitude higher than the pre-event intensity levels. Slight enhancements of energetic particle intensities at \stereob{} were associated with another unrelated solar event. The release time of the particles at L1-observers determined with the VDA method corresponded to the time when the magnetic field line of L1-observers initially connected to the shock. We suggest that the L1-observers was directly connected to the portion of the shock with enough particle acceleration efficiency via the magnetic field line. We give a lower limit of the particle release time for \stereoa{} based on the TSA method. The particle release of \stereoa{} was later than the release of the particles at L1-observers. It may be due to the fact that the magnetic connectivity between \stereoa{} and the portion of the shock with enough particle acceleration efficiency occurred later. Determining the arrival of EUV waves at the magnetic footpoints of the spacecraft is essential to test whether the EUV wave is responsible for the particle release. The EUV observations show that coronal holes and active regions influenced the propagation of the EUV wave, and no definite EUV wave arrived at the estimated magnetic footpoints of the spacecraft. We track the evolution of the EUV wave along the paths from the active region to the magnetic footpoints of L1-observers and \stereoa. It shows that the propagation of the EUV wave toward the magnetic footpoint of L1-observers was halted at about 18:35 UT. The EUV wave may be restrained by the coronal hole around the magnetic footpoint of L1-observers. The signature of the EUV wave toward the magnetic footpoint of \stereoa{} became diffuse and disappeared before the release of the particles at \stereoa. As a result, we do not obtain any definite evidence in the current case to support the argument that the particle release is accounted for by the arrival of the EUV wave at the magnetic footpoint of the spacecraft. This is consistent with previous studies of \citet{Gomez-Herrero15} and \citet{Lario14, Lario16}. However, we cannot exclude the possibility that EUV waves contribute to the particle releases in some cases. They may be able to explain the release of the first particles in those cases, but not the continuous releases which should be produced by the connection between the shock with enough particle acceleration efficiency and the observers. The geometrical modeling of the shock gives the 3D shock structure and allows us to analyze its contact with the magnetic field lines of the spacecraft. \citet{Lario16} suggest a relatively high altitude (${\gtrsim}$2 R$_{\odot}$ above the solar surface) for the height of the cobpoint when the particle were released. The fitting result in the current case gives relative lower and higher altitudes, i.e., a heliocentric height of 2.4${\pm}$0.5 R$_{\odot}$ for the L1 cobpoint when the particles were released and of 8.2${\pm}$0.7 R$_{\odot}$ for \stereoa{} cobpoint at the lower limit of the particle release time. The time when the magnetic field line of L1-observers initially connected to the shock was in good agreement with the release time of the particles observed at L1. In contrast, the release time of the particles at \stereoa{} was later than the time when the magnetic connectivity between \stereoa{} and the shock was first established. We suggest that \stereoa{} was first connected to a wave propagating backward toward the Sun via the magnetic field line, not the part of the shock with enough particle acceleration efficiency \citep{Liu17}. Particle acceleration efficiency of the shock is correlated with the parameters including the density compression ratio, the angle between the shock normal and the upstream magnetic field, the Mach number and the position angle relative to the shock leading edge \citep[e.g.,][]{Bemporad11, Kwon17, Kwon18, Lario17_Mach}. The fitting result shows $\theta_{Bn}$ of $\sim$51\degree{} and $\sim$31\degree{} at the cobpoints of L1-observers and \stereoa{} when the particles were released. The Mach number evaluated by the empirical models gives a modest value of ${\sim}$1.5 for the cobpoints of both L1-observers and \stereoa{} at the particle release times. It is comparable to the critical Mach number provided by \citet{Edmiston84}. We suggest that the particle acceleration efficiency of the portion of the shock connected to the spacecraft determines the release of energetic particles at the spacecraft. We confirm that the propagation of the EUV wave is affected by coronal holes and other active regions and it is not responsible for the particle release in our case. A possible role of particle perpendicular diffusion may contribute to the release of particles observed at the spacecraft, but this is beyond the scope of this paper. | 18 | 8 | 1808.04934 |
1808 | 1808.06931_arXiv.txt | { Chamaeleon I represents an ideal laboratory to study the cluster formation in a low-mass environment. % Recently, two sub-clusters spatially located in the northern and southern parts of Chamaeleon I were found with different ages and radial velocities.} {In this letter we report new insights into the structural % properties, age, and distance of Chamaeleon I based on the astrometric parameters from {\it Gaia} data-release 2 (DR2). % } {% We identified 140 sources with a reliable counterpart in the {\it Gaia} DR2 archive. We determined the median distance of the cluster using {\it Gaia} parallaxes % and fitted the distribution of parallaxes and proper motions assuming the presence of two clusters. % We derived the probability of each single source of belonging to the northern or southern sub-clusters, and compared the HR diagram of the most probable members % to pre-main sequences isochrones. % } {The median distance of Chamaeleon I % is $\sim$190\,pc. This is % about 20 pc larger than the value commonly adopted in the literature. From a {\it Kolmogorov-Smirnov test} of the parallaxes and proper-motion distributions % we conclude that the northern and southern clusters do not belong to the same parent population. % The northern population has a distance d$_{\rm N}$ = 192.7$^{+ 0.4}_{- 0.4}$ pc, while the southern one d$_{\rm S}$ = 186.5$^{+ 0.7}_{- 0.7}$ pc. The two sub-clusters appear coeval, at variance with literature results, and % most of the sources are younger than 3 Myr. The northern cluster is more elongated and extends towards the southern direction partially overlapping with the more compact cluster located in the south. A hint of a relative rotation between the two sub-clusters is also found. } {} | Chamaeleon I is one of the closest low-mass star forming regions with which it is possible to study all the key processes related to the formation of a young cluster, as well as the structure of protoplanetary disks around young stellar objects. \\ The stellar population of Chamaeleon I has been extensively investigated in the last fifteen years \citep[][]{FeigelsonLawson2004, Stelzeretal2004, Comeronetal2004}. The cluster is composed of two sub-structures, one northern and one southern. A complete study was presented by \citet{Luhman2007} who found the two populations to have different ages: 5-6 Myr for the northern and 3-4 Myr for the southern sub-cluster, respectively. The structure and dynamical properties of Chamaeleon I have been deeply investigated by \citet{Saccoetal2017}, % who % confirmed the presence of the two sub-clusters kinematically, with a shift in velocity of about 1 km/s. % \noindent The literature value of the distance of Chamaeleon I commonly adopted is 160 $\pm$ 15 pc \citep{Whittetetal1997}. This value comes from the combination of studies that employed different techniques. In particular, the extinction analysis of % \citet{Whittetetal1997} constrained the distance in the range between 135 and 165 pc, while the weighted average of Hipparcos distances for three cluster members \footnote{\object{HD 97300}, \object{HD 97048} and \object{CR Cha}}, is 175$^{+20}_{-16}$ pc \citep{perrymanetal1997}. \\%Combining this distance and the constraint of 135 - 165 pc {\it Gaia} DR2 astrometry \citep[][]{GaiaCollaboration2018, Lindegrenetal2018} clearly offers a unique opportunity to gather a new view of the region. \noindent In this letter we present the parallaxes and proper motions of the cluster members spectroscopically identified by previous works. The analysis of the {\it Gaia} DR2 data and % the discussion of the clusters kinematics and age are in Sects. \ref{an1} and \ref{disc}, respectively. | \label{disc} In this section we discuss the age and the kinematic properties of the north and south sub-clusters of Chamaeleon I. \noindent \subsection{The age of Chamaeleon I} In order to investigate whether an age difference is present between the two sub-clusters, we consider the 107 members with a probability higher than 80$\%$ defined in Section~\ref{an1}. We use the log(T$_{\rm eff}$)-M$_{\rm J}$ diagram in order to minimize the effects due to infrared excesses caused by the presence of protostellar disks. The effective temperatures are compiled either from \citet{Luhman2007} or from \citet{Saccoetal2017}. The absolute J magnitude of each source has been derived by adopting the mean distance module for the northern and southern sub-clusters and correcting for $A_J$ \citep[from][]{Luhman2007}. The overplotted isochrones are the Z=0.013 models from \citet{Tognellietal2011} \footnote{\url{https://www.astro.ex.ac.uk/people/timn/tau-squared/pisa_details.html}}. These models have a solar metallicity, which is a good approximation for the metallicity of the cluster as found by \citet{Spinaetal2017}.\\ \noindent As shown in Figure~\ref{prob}, all the sources have ages lower than 5 Myr. In particular, apart for a few sources, most of the members are younger than 3 Myr. While \citet{Luhman2007} found different ages for the two populations (5-6 Myr for the northern one and 3-4 Myr for the southern one), we do not find any evidence of an age difference between the two sub-clusters. \\ \noindent Our new findings and the differences with respect to the \citet{Luhman2007} results can be ascribed % to two effects: on one hand, \citet{Luhman2007} used the same distance to all sources, adopting a smaller value than what we find here; on the other hand, his selection of the two sub-clusters was based only on the spatial distribution, while in our case we take into account also their different parallaxes and kinematic properties. \subsection{Kinematics properties of the north and south sub-clusters} Under the assumption of an isotropic distribution in a star cluster, we can use the relation of \citet{platais_2012} to derive the velocity dispersion from the proper motion dispersion: \begin{equation} \sigma_r(km/s)=d\,(kpc)\,\cdot\,4.37\,\cdot\,\sigma_\mu(mas/yr) ,\end{equation} where the $\sigma_\mu^2 = \frac{\sigma_{\mu_\alpha}^2+\sigma_{\mu_\delta}^2}{2}$ \citep[][]{McLaughlinetal2006}. \\ \noindent We obtain $\sigma_{r,N}= 0.681\pm 0.057$ km/s and $\sigma_{r,S}= 0.727\pm 0.134$ km/s, where the uncertainties are computed from the error propagation. The velocity dispersions are consistent, within 2 $\sigma$, with the results of \citet{Saccoetal2017}. % \noindent Given that the northern cluster is in the background, and it is more redshifted than the closer southern sub-cluster, we conclude that the two clusters are moving away from each other. \noindent In Figure~\ref{histplx2} the two arrows represent the proper motions of the two sub-clusters with respect to a reference system centered on the cluster. This confirms that the two sub-clusters are not merging and have a non-zero angular momentum. Combining this result with the differential radial velocity measured by \citet{Saccoetal2017}, this represents a hint of rotation of the two sub-clusters. This is a new and puzzling result. Indeed, in young high-mass clusters rotation has been theoretically predicted by \citet{Mapelli2017}, and it has been observed, for example, in the high-mass star forming region R136 in the Large Magellanic Cloud \citep[][]{Henault-Brunetetal2012}. However, in simulations with similar total mass to low-mass environments, such as Chamaeleon I, \citet{Mapelli2017} did not find a clear signature of rotation as in high-mass environments. \\%Our simulated rotation curves \noindent \begin{figure*}% \begin{minipage}[b]{0.5\textwidth} \includegraphics[trim=0cm 4cm 0cm 0cm,width=0.9\textwidth]{fig/mle_hist_plx1N.jpg} \\ \includegraphics[width=0.9\textwidth]{fig/mle_hist_plx1S.jpg} \end{minipage} \begin{minipage}[b]{0.5\textwidth} \includegraphics[trim=15cm 6cm 0cm 0cm,width=1.1\textwidth]{fig/pos_pb.jpg} \end{minipage} \caption{{\it Left:} Parallax distribution of the most probable members (P$\ge$80$\%$) of the northern (in blue) and southern (in magenta) sub-clusters. In each panel the MLE parallaxes for both clusters are shown. % {\it Right:} Spatial distribution of the northern and the southern sub-cluster (as in the {\it left} panel). The red arrows represent the differential proper motions in $\alpha$ and $\delta$ with respect to a mean proper motion between the northern and southern cluster. } \label{histplx2} \end{figure*} \begin{figure*}[htb] \centering% \includegraphics[trim=5cm 3cm 1cm 4cm,width=13cm]{fig/mle_cmdkTeff_new.jpg} \caption{log(T$_{\rm eff}$) - M$_{\rm J}$ Diagram of the most probable members (P$\ge$80$\%$). The color code is as in Figure \ref{histplx2}. The solid lines are the pre-main sequence isochrones at different ages between 1 and 10 Myr. .} \label{prob} \end{figure*} A more detailed analysis of the cluster dynamics will be presented in an upcoming paper, % together with an updated census of the members using {\it Gaia} DR2 data. | 18 | 8 | 1808.06931 |
1808 | 1808.00330_arXiv.txt | { The discovery of the Fermi bubbles---a huge bilobular structure seen in GeV gamma-rays above and below the Galactic center---implies the presence of a large reservoir of high energy particles at $\sim 10 \, \text{kpc}$ from the disk. The absence of evidence for a strong shock coinciding with the edge of the bubbles, and constraints from multi-wavelength observations point towards stochastic acceleration by turbulence as a likely mechanism of acceleration. We have investigated the time-dependent acceleration of electrons in a large-scale outflow from the Galactic centre. For the first time, we present a detailed numerical solution of the particle kinetic equation that includes the acceleration, transport and relevant energy loss processes. We also take into account the addition of shock acceleration of electrons at the bubble's blast wave. Fitting to the observed spectrum and surface brightness distribution of the bubbles allows determining the transport coefficients, thereby shedding light on the origin of the Fermi bubbles. } | The detection of the \textit{Fermi} bubbles--a huge bi-lobular structure seen in GeV gamma-rays--is certainly one of the great discoveries made with the \textit{Fermi}-LAT instrument. Due to their position on the sky (see below), they are likely emanating from the Galactic centre and the most speculated about sources are the supermassive black hole at the Galactic centre and star formation/star burst in the Galactic centre region. These processes shape Galactic structure on the largest scales and as such the Fermi bubbles allow us to study Galactic feedback in our own backyard. Furthermore, given their prominence in gamma-rays, they are an important arena for studies of sources of diffuse GeV emissions, like searches for signals from self-annihilation or decay of dark matter. Finally, the production of the gamma-rays and the acceleration of the underlying particles are of astrophysical interest in itself. \subsection{Observational properties} Originally the Fermi bubbles were observed in a search \citep{Dobler:2009xz} for the gamma-ray counterpart of a microwave excess seen from the inner Galaxy \citep{Finkbeiner:2003im,Dobler:2007wv,Ade:2012nxf}. A more detailed analysis \citep{Su:2010qj} unveiled some surprising properties that were later largely confirmed by \citet{Fermi-LAT:2014sfa}. In the following we summarise the most important observational properties of the Fermi bubbles in gamma-rays. \paragraph{Geometry:} The Fermi bubbles are approximately centered at zero Galactic longitude, symmetric about the Galactic plane, $50^\circ$ wide in longitude with each bubble extending up to $50^\circ$ in latitude, see, e.g.\ Fig.~22 of \citet{Fermi-LAT:2014sfa}. On these scales, they constitute the first evidence for an outflow from the Milky Way. (On smaller scales, there had previously been evidence in X-rays in an X-shaped feature around the Galactic centre.) The bubbles' symmetry about the Galactic plane and their being centred around zero longitude imply an origin at Galactic centre (distance $d_{GC} \simeq 8.5 \, \text{kpc}$). A wind with a constant speed of $1000 \, \text{km} \, \text{s}^{-1}$ would need about $9.9 \, \text{Myr}$ to expand into a bubble of size $\sim d_{GC} \tan 50^{\circ} \simeq 10.1 \, \text{kpc}$, modulus projection effects: At a latitude of $50^{\circ}$, we might be seeing the limb-brightened edge of a bubble of radius $d_{GC} \sin 50^{\circ} \simeq 6.5 \, \text{kpc}$, thus reducing the time-scale to $6.4 \, \text{Myr}$. Note that because the eastern edge of the northern bubble is very close to the position of the North-polar spur, which is part of the the radio Loop~I. Initially, this led to claims of the bubbles being associated with the Loop~I structure \citep{Casandjian:2009wq}. \paragraph{Spectrum:} The gamma-ray flux shows a hard spectrum, mostly $\propto E^{-2}$ and extending from a few hundred MeV up to a few hundred GeV, see, e.g.\ Fig.~18 of \citet{Fermi-LAT:2014sfa}. At lower energies, the spectrum is significantly harder, and at high energies there is evidence for a spectral softening or an exponential cut-off. This spectral shape immediately invites speculation about its physical origin, i.e. whether the gamma-rays are of leptonic (from inverse-Compton scattering) or hadronic ($\pi^0$ decay) origin. (Given the estimates of the physical conditions, inside the bubbles, see below, bremsstrahlung is most likely negligible.) While the spectral shoulder around a few hundred MeV determined in the earlier analysis \citep{Su:2010qj} seemed to be well fit by the kinematic feature from $\pi^0$ decay, the new best-fit spectrum appears to be extending to lower energies. Likely, a hadronic model needs to have a spectral break (a steeper spectrum of the underlying protons) at lower energies. This is in addition to the required spectral break or cut-off at high energies. The physical origin of these breaks is \textit{a priori} unclear. In leptonic models these breaks are easily explained. The inverse-Compton spectrum is naturally rather hard: In the Thomson regime, a gamma-ray spectrum $\propto \epsilon^{-s}$ with ($s\sim 2)$ with a cutoff at $\epsilon_{\rm cut}$ of a few hundred GeV can be produced by an electron spectrum $\propto E^{-\Gamma} \mathrm{e}^{-E/E_{\mathrm{cut}}}$ with $\Gamma = 2s-1\sim 3$ and $E_{\mathrm{cut}}=/m_{\mathrm{e}} c^2\sqrt{\epsilon/\epsilon_{\rm soft}}\sim 1,500$ GeV for soft photon energies of $\epsilon\sim 1$ eV. Note that these estimates are strictly only valid in the Thomson regime. In the numerical computations, however, we have used the full Klein-Nishina cross-section and taken into account the relativistic corrections. \paragraph{Surface brightness.} The surface brightness shows little variation over the bubbles, but has sharp edges as can already be seen in the residual map, cf.~e.g.~Fig.~29 of \citet{Fermi-LAT:2014sfa}. More quantitatively, this is evidenced by profiles of the gamma-ray flux across the bubble edge, shown e.g. in Fig. 22 of \citet{Fermi-LAT:2014sfa}. There is clearly a jump in intensity from a value close to zero (after template subtraction) outside to a relatively constant value inside the bubbles. In fact, the only substructure seen is a rather large enhancement of emissivity in the east of the southern bubble, called the ``cocoon'', the origin of which is of yet unknown. There have also been claims of evidence for a narrow and extendend, jet-like feature \citep{Su:2012gu}, however, the analysis by the Fermi collaboration \citep{Fermi-LAT:2014sfa} has found this feature not to be significant. The flat surface brightness and sharp edges are one of the most puzzling features of the bubbles. The sharp edges require an efficient confinement of the gamma-ray producing particles and the flat surface brightness requires a peculiar distribution of volume emissivity. Idealising each bubble as a spherically symmetric volume with outer radius $R$, only an emissivity that varies with radius $r$ as $1 / \sqrt{1 - (r/R)^2}$ will give a flat surface brightness and sharp edges. \paragraph{Spectral uniformity.} The bubbles show similar morphologies in different energy bins ranging from $1$ to $500 \, \text{GeV}$ (see e.g.~Fig.~22 of \citet{Fermi-LAT:2014sfa}) or equivalently the spectrum is uniform in different parts of the bubbles. Specifically, the gamma-ray spectrum has been analyzed in different latitude bands and the spectrum in the bubble edge region and the interior have been compared: For the latitude bands, no variation has been found above and below $\pm 10^{\circ}$. Between $-10^{\circ}$ and $+10^{\circ}$ there is an excess at the Galactic Centre \citep{Hooper:2013rwa}, likely with a spherical symmetry, and its connection to the Fermi bubbles is unclear at this point \citep{TheFermi-LAT:2017vmf}. Furthermore, no variation between the edge region and the interior was found \citep{Su:2010qj} (but see also \citet{Keshet:2016fbq}). The spectral uniformity is also very surprising for such an extended structure. Leptonic models in particular would be expected to lead to some variation, depending on the region of energizing of the high-energy electrons. This is due to cooling losses by synchrotron radiation and inverse-Compton emission. A conservative estimate of the cooling time is $\tau_{\text{cool}} = 6.74 \times 10^8 \text{yr} \, (E/\text{GeV})^{-1} ((u_B + u_{\text{CMB}})/ (0.486 \, \text{eV} \text{cm}^{-3}) )^{-1}$, for magnetic fields and radiation fields of energy densities $u_B = 0.224 \, \text{eV} \, \text{cm}^{-3} (B / 3 \mu\text{G})^2$ and $u_{\text{CMB}} = 0.262 \, \text{eV} \, \text{cm}^{-3}$, respectively, i.e. of the same order as the bubble age for $100 \, \text{GeV}$ electrons. Therefore, electrons energised in the Galactic plane will be subject to considerable cooling while travelling out into the bubble volume. This results in softer spectra at larger distances from the Galactic centre and thus a softer gamma-ray spectrum at higher latitudes. In addition, the energy densities in the radiation backgrounds that the electron inverse-Compton scatter on should be varying with distance from the disk: While the CMB is of course spatially uniform, the energy densities in both the optical/UV and the infrared backgrounds should become smaller further away from the disk. The fact that this is not observed imply that the variation in the radiation backgrounds must be counter-balanced by a variation in the electron spectrum to some degree. \subsection{Hints} While the discovery of the Fermi bubbles was certainly a surprise, it was not the first hint at the presence of Galaxy-scale outflows. Kiloparsec-scale outflows have been observed for starburst galaxies, e.g. in ionised gas. Even in our own Milky Way, there had been hints at the presence of a Galaxy-scale outflow, possibly connected with high-energy cosmic rays: Observations in soft X-rays, most notably from ROSAT, showed signs of an x-shaped feature, interpreted as evidence of a biconical outflow in analogy with structures seen in other galaxies. The presence of a population of high-energy cosmic ray electrons was already hinted at by the microwave haze, an excess of microwaves from the Galactic centre, pointing at a similarly hard electron spectrum \citep{Finkbeiner:2003im,Dobler:2007wv,Ade:2012nxf}. (Note, however, the possible influence of systematic effects due to template subtraction \citep{Mertsch:2010ga}.) The search for a counterpart of the microwave haze in gamma-rays was in fact what motivated the first study that lead to the discovery of the Fermi bubbles~\citep{Dobler:2009xz}. \subsection{Other constraints} \label{sec:other_constraints} \paragraph{X-rays.} A number of studies have investigated the properties of the thermal gas in the Fermi bubbles and in the Galactic halo from X-ray observations. The parameters can be either inferred from the thermal, soft X-ray spectrum \citep{Kataoka:2013tma,Kataoka:2015dla} or from individual Oxygen lines \citep{Miller:2016chr}. The gas densities inferred are of the order $n_{\text{gas}} \sim 10^{-3} \, \text{cm}^{-3}$ and the temperatures of the gas just outside the bubbles vary between $kT \simeq 0.3 \, \text{keV}$ and $0.5 \, \text{keV}$. This is higher than the canonical temperature of the Galactic halo of $kT \simeq 0.2 \, \text{keV}$ and requires a heating agent, perhaps a weak shock with a low Mach number; $\mathcal{M} \simeq 1.5 \mathellipsis 2.3$. Finally, with the typical sound speed in the Galactic halo of $c_s \simeq 200 \, \text{km} \, \text{s}^{-1}$, one infers shock speeds of $v_{\rm sh} \simeq 300 \mathellipsis 500 \, \text{km} \, \text{s}^{-1}$. The absence of evidence for a strong shock coinciding with the bubble edge implies that diffusive shock acceleration at the bubble edge cannot be the primary mechanism of acceleration. If electrons get accelerated in the Galactic plane or even in a hypothetical large-scale jet along the Galactic minor axis, they need to travel over distances of several kpc without much energy loss to fill the bubble volume. As a result they will suffer severe cooling losses and a gradual softening of their spectrum, or even quench the electron density completely. Note further that the low shock speeds found by the X-ray modeling lead to even larger dynamical times than with the $1000 \, \text{km} \, \text{s}^{-1}$ assumed above, making the energy losses even more important. \paragraph{Quasar absorption.} The observation of absorption by the gas associated with the bubbles from a background quasar can also be used to set bounds on the outflow speed. In the UV absorption lines from PDS 456 two (asymmetric) components with velocities of $v \simeq -235$ and $+250 \, \text{km} \, \text{s}^{-1}$ with respect to the local standard of rest could be identified \citep{Miller:2016chr}. For the conical outflow assumed in that study, this implies an upper limit on the outflow speed of $\gtrsim 900 \, \text{km} \, \text{s}^{-1}$. This seems to be in conflict with the shock speed inferred from the X-ray modelling described above. Note, however, that the outflow speed inferred from the absorption lines of one quasar is very dependent on the assumed geometry of the flow. Future observations of additional sight lines towards other quasars can help mapping out the flow structure, thus possibly also constraining it geometry, and might bring the results into agreement with the values inferred from X-rays. \subsection{Models} \label{sec:models} The Fermi bubbles have also generated a great deal of interest on the modelling side \citep{Crocker:2010dg,Cheng:2011xd,Cheng:2011tx,Zubovas:2011py,Mertsch:2011es,Zubovas:2012bn,Guo:2011eg,Yang:2012fy,Lacki:2013zsa,Crocker:2013mna,Fujita:2013jda,Thoudam:2013eaa,Yang:2013kca,Crocker:2013mna,Fujita:2014oda,Cheng:2014nva,Cheng:2014lca,Mou:2014pea,Crocker:2014fla,Cheng:2015zda,Mou:2015wxa,Sarkar:2015xta,Sasaki:2015nta,Yang:2017tjr}. The variety of models is most conveniently classified by: \begin{itemize} \item the source of energy: super massive black hole or stellar winds/supernovae; \item the acceleration region: jet or sources in the disk or \textit{in situ} (by shocks or turbulence); \item the nature of the high-energy particles: hadrons or leptons. \end{itemize} Of course, the individual options are not mutually exclusive. For instance, in hadronic models, the bulk of the high-energy gamma-rays comes from decay of neutral pions. Charged pions, however, get produced at similar rates and can, given the radiation fields, their $e\pm$ byproducts can inverse-Compton scatter soft photons into low energy gamma-rays. However, in this particular scenario the synchrotron spectrum would be too soft (Ackermann et al. 2014) As a full discussion of all proposed models is beyond the scope of this highlight presentation, only two particular classes of models will be presented, and a few concrete examples will be shown. \paragraph{Jet models.} Astrophysical jets are thought to be powered by accretion onto a spinning, compact object, like neutron stars or black holes. Given the position and symmetry of the Fermi bubbles, the supermassive black hole at the Galactic centre is a prime candidate. Although conspicuously quiet (its X-ray luminosity is currently more than 11 orders below the Eddington luminosity), there is indirect evidence for earlier epochs of active accretion, e.g. from X-ray reflections. Jets are usually associated with high speeds $\gtrsim 1000 \, \text{km} \, \text{s}^{-1}$. This allows for the electrons to be less impacted by energy losses than in starburst/star formation models and therefore the source of energisation of the high-energy electrons can be in the Galactic disk or inside the jet. (Note, however, that the jet speed is not necessarily directly implying the dynamical age as the bubbles can be formed by a fountain-like back flow due to the termination of the jet by the ram pressure of gas in the Galactic halo.) One of the earliest studies of a leptonic jet model employing a hydrodynamical simulation \citep{Guo:2011eg} found that the lateral extent of the Fermi bubbles could be explained if the jet was underdense but slightly overpressured. If active at $10 \, \%$ of the Eddington luminosity for $1 - 2 \, \text{Myr}$ until about a $\text{Myr}$ ago, the morphology would match the observations. A subsequent MHD simulation of the Fermi bubbles blown up by a jet \citep{Yang:2012fy} showed further that the shock compression at the bubble edges would compress the magnetic field such that it gets aligned with the bubble edge. We will return to this point in Sec.~\ref{sec:toroidal_coordinates}. \paragraph{Star formation/star burst models.} The Galactic winds that get collectively powered by an ensemble of stellar winds or supernova activity, are usually operating at smaller speeds, $\lesssim 500 \, \text{km} \, \text{s}^{-1}$. This implies a larger dynamic time-scales than for the jet model, leading to a preference in the literature for hadronic models, as leptons would cool too fast. In hadronic models, on the other hand, cosmic rays need to be accumulated over much longer time scales, given the low gas densities of the order of $10^{-3} \, \text{cm}^{-3}$ (see Sec.~\ref{sec:other_constraints}), to produce the observed gamma-ray fluxes. In turn, this and the observed hard $E^{-2}$ spectrum require an effective confinement of the high-energy cosmic rays to the bubbles and a suppression of (energy-dependent) escape. (See, however, \citet{Keshet:2016fbq}.) The sources of high energy particles are nevertheless oftentimes assumed to be in the Galactic disk. The most detailed numerical star formation/star burst model for the Fermi bubbles as of yet \citep{Sarkar:2015xta} employs a hydrodynamical code to investigate the interaction of a Galactic wind with the circumgalactic medium. It is found that a luminosity of $5 \times 10^{40} \, \text{erg} \, \text{s}^{-1}$ and a density in the halo of $10^{-3} \, \text{cm}^{-3}$ can reproduce the morphology observed in gamma-rays and is also in agreement with X-ray observations. Interestingly, this luminosity is close to the one inferred from the current star formation rate, $\text{SFR} \simeq 0.007 \mathellipsis 0.1 M_{\odot} \, \text{yr}^{-1}$, when assuming an efficiency of $30 \, \%$ for conversion into mechanical power, $\mathcal{L} \simeq 10^{40} \text{erg} \, \text{s}^{-1} \varepsilon_{0.3} (\text{SFR} / (0.1 M_{\odot} \, \text{yr}^{-1}))$. The outflow from the inner Galaxy leads to a shock structure known from the heliosphere or supernova remnants, with a radial forward shock at $\sim 11 \, \text{kpc}$, a more tangled contact discontinuity extending to $\sim 8 \, \text{kpc}$ above the Galactic centre and a very much tangled reverse shock a few kiloparsecs inside of the contact discontinuity. Thus, in this model, the edge of the gamma-ray bubble does not coincide with the projection of the forward shock, but rather the contact discontinuity. Whether this is due to the diffusion prescription of \citet{Sarkar:2015xta} changing across the contact discontinuity would need to be explored further. \subsection{Motivation} The observation of $\gamma$-rays from the bubbles implies a huge reservoir of high-energy particles in the Galactic halo, but the source and the mechanism of acceleration of these particles has not been established thus far. Other sources of non-thermal particles, like supernova remnants, pulsar wind nebulae, jets in active galaxies or winds in starburst galaxies, show evidence of shocks through X-rays or ionization lines. The Fermi bubbles, however, show no such evidence of a (strong) shock, raising the question of the possible mechanism of acceleration. Acceleration by plasma turbulence (or ``stochastic acceleration'', SA), however, can fill the bubbles with high-energy electrons. (See, \citet{Petrosian:2012ba} for a recent review of SA. A first SA model for the Fermi bubbles \citep{Mertsch:2011es} was presented quickly after their discovery. This model was employing the solution of a simplified version of the transport equation. Specifically, diffusion was ignored as a spatial transport process and advection was the only transport process. In this framework, cosmic ray electrons are just passively advected with the downstream flow while being stochastically accelerated. The time scale hierarchy $t_{\text{dyn}} \gg t_{\text{cool}} \gg t_{\text{acc}}$ of dynamical, cooling and acceleration times, allows a steady-state solution of the variation of electron spectrum with radius for a given spatial variations of theses and and the escape time, $t_{\text{esc}}$. While successful in explaining the overall spectrum of the bubbles as well as the sharp edges, the lack of diffusive transport was an important shortcoming. In addition, the interstellar radiation fields on which the cosmic ray electrons scatter was assumed homogeneous which must be an oversimplification. What is needed is a detailed numerical model, taking into account all the spatial transport processes (diffusion, advection), energy losses (ionisation, bremsstrahlung, synchrotron, inverse Compton scattering) and energy gains (shock and SA). In the remainder of this paper, we will present our computation of the SA of high-energy electrons in the Fermi bubbles. Sec.~\ref{sec:method} introduces our method for solving the transport equation on a grid that is suited for the geometry of the bubbles. We will define three setups and specify the parameter values considered. We will show our result for those three setups in Sec.~\ref{sec:results} and comment on compatibility with observational data. In Sec.~\ref{sec:summary_conclusions} we summarise and conclude. | \label{sec:summary_conclusions} \begin{enumerate} \item We have reviewed the observations of the Fermi bubbles in gamma-rays relevant for the modeling of the emission, transport and acceleration processes with particular focus on their puzzling morphological and spectral properties; namely the constancy of their surface brightness with abrupt edges and relatively uniform hard spectra. We discussed X-ray emission and UV absorption line observations which give (somewhat conflicting) bounds on the outflow velocities in the bubbles and include in our discussion the observations of so-called microwave haze. We reviewed briefly some of the models proposed for production of the bubbles including MHD simulation, jet and Galactic wind models. \item Our main focus, however, is the acceleration of particles, their transport and emission characteristics with the primary goal of explaining the puzzling morphological and spectral characteristics. We present arguments in favor of leptonic vs.\ hadronic model and develop kinetic equations to described the acceleration and transport of electrons throughout the bubbles. We include effects of stochastic acceleration by turbulence and those of a low Mach number shock with special attention to momentum and spatial diffusion coefficient in the magnetized medium of the bubble. We also include the effects of energy loss due to inverse Compton and synchrotron processes in a inhomogeneous magnetic and soft photon (CMB, infrared and optical/UV) fields. \item We present results from three different models with similar characteristics of the (better understood) loss mechanisms but with different assumptions about more uncertain acceleration and other transport characteristics. The first model has isotropic spatial diffusion in the bubble ($\eta=1$) and anisotropic ($K_{uu} \ll K_{vv}$; $\eta=10$) in the halo and higher acceleration rate inside the bubble. This model results in surface brightness distribution not as uniform as observed and can be ruled out. The second model has anisotropic diffusion both inside ($\eta=3.16$) and (even stronger) outside ($\eta=31.6$) in the halo. This model results in a uniform surface brightness but the bubble edge is not as sharp as observed. The third model consists of a relatively thick finite size shell expanding into the halo with mildly anisotropic diffusion in the shell ($\eta=3.16$) but with a much stronger anisotropy inside and outside in the halo ($\eta=31.6$). Acceleration rate is ten times higher in the shell than outside. This model produces a sharper edge and agrees with the observed spectral distribution in gamma-rays and microwave ranges. \item We conclude that the gamma-ray as well as microwave spectral and morphological features of the Fermi bubbles can be reproduced by the Inverse Compton and synchrotron emission from electrons accelerated by turbulence generated in a mildly supersonic outward flowing shell. This finding is strengthening the scenario where the bubbles are inflated by a wind powered by star formation or star burst activity. Another possibility for inflating the bubbles is a jet from past AGN activity at the center of the galaxy. Whether an \textit{in situ} acceleration of particles in the jet environment can lead to explanation of observed characteristics of the bubbles as done by our model would require a separate study. If such a future study were to conclude that jet models could not produce the observed properties, it would strengthen the above conclusions based on our current study.\end{enumerate} | 18 | 8 | 1808.00330 |
1808 | 1808.09816_arXiv.txt | {The \textit{INTEGRAL} satellite has been observing the $\gamma$-ray sky for 15 years and has detected over 900 X-ray sources of various nature. However, more than 200 of these sources still lack precise identification.} {Our goal is to reveal the nature of the high-energy sources detected by \textit{INTEGRAL}. In particular, we want to improve the census of X-ray binaries.} {Photometry and spectroscopy were performed in July 2012 on 14 \textit{INTEGRAL} sources in near-infrared at the Very Large Telescope on the European Southern Observatory-UT3 telescope equipped with the ISAAC spectrograph. We used K$_s$ images reaching to a depth of magnitude 18.5 to look for unique counterparts to high-energy detections to check for both extended sources and photometric variability. The analysis of near-infrared spectral features allows us {to constrain} the nature of these X-ray sources by comparing them to {stellar} spectra atlases.} {We present photometric and/or spectroscopic data for 14 sources (IGR J00465$-$4005, IGR J10447$-$6027, IGR J12489$-$6243, IGR J13020$-$6359, IGR J13186$-$6257, IGR J15293$-$5609, IGR J17200$-$3116, IGR J17404$-$3655, IGR J17586$-$2129, IGR J17597$-$2201, IGR J18457+0244, IGR J18532+0416, IGR J19308$+$0530, and IGR J19378$-$0617). We conclude that 5 of these are active galactic nuclei, 5 are cataclysmic variables, 2 are low- or intermediate-mass X-ray binaries, and 2 are Be high-mass X-ray binaries.} {} | \label{sect:intro} The \textit{INTEGRAL} satellite has been observing the high-energy sky between 15 keV and 10 MeV for 15 years. However, the nature of the sources detected at high energies is often uncertain and requires further observations at low energies in optical and near-infrared (nIR) wavelengths. This is why a significant fraction ($\sim$\,20\%) of the \textit{INTEGRAL} detections still need better constraints to have a robust identification. According to the catalogue of \textit{INTEGRAL} detections published by \cite{bird_ibis_2016}, it is expected that a majority of the unknown sources are active galactic nuclei (AGN) and high-energy binaries; the latter are binary star systems that host an accreting compact object. High-energy binaries fall into three main categories depending on the compact object and the mass of the companion star: cataclysmic variables (CV), low-mass X-ray binaries (LMXB), and high-mass X-ray binaries (HMXB). While LMXBs host either a neutron star (NS) or a black hole (BH), CVs host a white dwarf\,; both have a low-mass companion star (typically M\,$\leq$\,1\,M$_\sun$). High-energy radiation is released through accretion of matter from the companion star overflowing its Roche lobe. An accretion disc can form around the compact object and feed it gradually, which often leads to transient behaviour. The less common intermediate-mass X-ray binaries (IMXBs) have a companion of mass between 1 and 10\,M$_\sun$ and the accretion process is similar to that of LMXBs. For the sake of consistency with the literature and especially \cite{bird_ibis_2016}, we group IMXBs with LMXBs. The HMXBs host a massive star (typically M\,$\geq$\,10\,M$_\sun$), around which orbits either a NS or a BH. Two main subcategories exist among HMXBs, based on the evolutionary phase of the companion star. In Be-types (BeHMXB), the secondary is a fast-rotating main-sequence O/B star that sheds matter from its equator as a consequence of high centrifugal force. A decretion disc thus forms around the companion star. The accretion of matter occurs when the compact object passes through that disc. Supergiant binaries \mbox{(sgHMXB)} host an evolved O/B supergiant star that feeds a compact object with matter through an intense stellar wind, driven by its tremendous luminosity. The sensitivity of \textit{INTEGRAL} at higher energies made it possible to differentiate two new subclasses of sgHMXBs, as reviewed in \cite{chaty_optical/infrared_2013}. Obscured HMXBs are intrinsically absorbed ($N_H > 10^{23}$\,cm$^{-2}$), while supergiant fast X-ray transients (SFXTs) show short and intense bursts of high-energy radiations. \begin{table*}[ht!] \centering \small \caption{Positions and uncertainty of the X-ray detections of the IGR sample.\label{tab:xpos}} \begin{tabular}{lllrrl} \hline\hline\\[-1.5ex] Source & RA J2000 & Dec J2000 & \textit{l} & \textit{b} & Uncertainty at 90\%\\ & (X-ray) & (X-ray) & (\degr) & (\degr) & (\arcsec) \\[1.5ex] \hline\\[-1.5ex] IGR J00465$-$4005 & 00:46:20.71 & $-$40:05:47.3 & 307.2621 & $-$76.9884 & 4\farcs26 (\textit{Swift})$\,^a$ \\ IGR J10447$-$6027 & 10:44:51.89 & $-$60:25:12.0 & 287.9185 & 1.2915 & 0\farcs65 (\textit{Chandra})$\,^b$ \\ IGR J12489$-$6243 & 12:48:46.44 & $-$62:37:43.1 & 302.6257 & 0.2417 & 0\farcs64 (\textit{Chandra}) $\,^c$\\ IGR J13020$-$6359 & 13:01:58.72 & $-$63:58:08.7 & 304.0885 & $-$1.1212 & 0\farcs39 (\textit{XMM}) $\,^d$ \\ IGR J13186$-$6257 & 13:18:25.08 & $-$62:58:15.5 & 305.9915 & $-$0.2599 & 0\farcs64 (\textit{Chandra}) $\,^e$ \\ IGR J15293$-$5609 & 15:29:29.37 & $-$56:12:13.3 & 323.6587 & 0.1712 & 0\farcs64 (\textit{Chandra}) $\,^c$ \\ IGR J17200$-$3116 & 17:20:05.92 & $-$31:16:59.4 & 355.0221 & 3.3472 & 0\farcs64 (\textit{Chandra}) $\,^f$ \\ IGR J17404$-$3655 & 17:40:26.86 & $-$36:55:37.4 & 352.6259 & $-$3.2725 & 0\farcs64 (\textit{Chandra}) $\,^e$ \\ IGR J17586$-$2129 & 17:58:34.56 & $-$21:23:21.6 & 7.9862 & 1.3265 & 0\farcs64 (\textit{Chandra}) $\,^e$ \\ IGR J17597$-$2201 & 17:59:45.52 & $-$22:01:39.17 & 7.5696 & 0.7704 & 0\farcs60 (\textit{Chandra}) $\,^g$ \\ IGR J18457$+$0244 & 18:45:40.38 & $+$02:42:09.2 & 34.6819 & 2.5135 & 2\farcs16 (\textit{XMM-Newton}) $\,^d$ \\ IGR J18532$+$0416 & 18:53:15.91 & $+$04:17:48.26 & 36.9651 & 1.5519 & 1\farcs14 (\textit{XMM-Newton}) $\,^d$ \\ IGR J19308$+$0530 & 19:30:50.77 & $+$05:30:58.09 & 42.3807 & $-$6.1852 & 0\farcs6~~ (\textit{Chandra})$\,^g$ \\ IGR J19378$-$0617 & 19:37:33.1 & $-$06:13:04 & 32.5905 & $-$13.0737 & 3\farcs5~~ (\textit{Swift}) $\,^h$ \\ \hline \end{tabular} \begin{minipage}{.77\textwidth} \begin{tiny} $^a$\cite{landi_swift/xrt_2010-1}, $\,^b$\cite{fiocchi_five_2010}, $\,^c$\cite{tomsick_localizing_2012}, $\,^d$\cite{rosen_XMM-Newton_2016}, $\,^e$\cite{tomsick_chandra_2009}, $\,^f$\cite{tomsick_chandra_2008}, $\,^g$\cite{ratti_chandra_2010}, $\,^h$\cite{rodriguez_swift_2008} \end{tiny} \end{minipage} \end{table*} Precisely identifying high-energy sources requires further observations, for which the nIR domain is well adapted. Firstly, many \textit{INTEGRAL} sources (later called IGRs) lie near the Galactic plane, where optical radiation is absorbed by dust while infrared is not. Secondly, most of a binary's nIR emission comes from the companion star or the accretion disc, which is ideal to identify their nature by constraining its spectral type. In this paper, we present a sample of 14 IGR sources (\tabref{tab:xpos}) for which we performed nIR photometry and/or spectroscopy. We aim to confirm a unique nIR counterpart for each of these IGRs and provide additional constraints on their nature such as the spectral type of companion stars in X-ray binaries. The IBIS instrument on board \textit{INTEGRAL} has a wide field of view, but does not have enough spatial resolution to associate accurately an optical/nIR counterpart to the high-energy detections. Precise X-ray localization is thus given by either \textit{Chandra}, \textit{XMM-Newton,} or \textit{Swift} telescopes. Section \ref{sect:observations} describes\ nIR photometric and spectroscopic observations by the European Southern Observatory (ESO), along with data reduction processes. Section \ref{sect:results} compiles all the previously published results on the sources along with the new results of our nIR observations. In Section \ref{sect:discuss} we discuss the results and outcomes of these observations before concluding in Section \ref{sect:conclusion}. | \label{sect:conclusion} We presented nIR observations of 14 \textit{INTEGRAL} sources with the VLT/ISAAC instrument. The photometric and spectroscopic data allowed us to pinpoint nIR counterparts to the high-energy detections and identify or better constrain the nature of the sources. Among these sources, there are 5 AGNs, 5 CVs, 2 BeHMXBs, and 2 I/LMXBs. While the proportions between types are not fully consistent with those published in \cite{bird_ibis_2016}, we still expect that the remaining unidentified \textit{INTEGRAL} sources contain a significant amount of AGNs, X-ray binaries, and CVs. This could be a great resource for the two latter, since the current census of binaries is not so high and would benefit from having more candidates with a well-constrained nature. In turn, this will help to perform population studies, derive accurate classifications, and answer more general questions on stellar evolution in binaries in the context of stellar merging endpoints and the detection of gravitational waves. | 18 | 8 | 1808.09816 |
1808 | 1808.00276_arXiv.txt | We conduct three-dimensional hydrodynamical simulations of two opposite jets that interact with a spherical slow wind that includes a denser shell embedded within it, and obtain a bipolar nebula where each of the two lobes is composed of two connected bubbles and Rayleigh-Taylor instability tongues that protrude from the outer bubble and form the `columns crown'. The jets are launched for a short time of 17 years and inflate a bipolar nebula inside a slow wind. When the bipolar structure encounters the dense shell, the interaction causes each of the two lobes to split to two connected bubbles. The interaction is prone to Rayleigh-Taylor instabilities that form tongues that protrude as columns from the outer bubble. The bases of the columns form a ring on the surface of the outer bubble, and the structure resemble a crown that we term the columns crown. This structure resembles, but is not identical to, the many filaments that protrude from the lobes of the bipolar planetary nebula Menzel~3. We discuss our results in comparison to the structure of Menzel~3 and the ways by which the discrepancies can be reconciled, and possibly turn our failure to reproduce the exact structure of Menzel~3 to a success with jets-shell interaction simulations that include more ingredients. | \label{sec:intro} The interaction of a binary companion with an evolve giant star, like an asymptotic giant branch (AGB) star that is a progenitor of a planetary nebula (PN), has in principle two sources of gravitational energy that might energize the outflow. The first one is the orbital gravitational energy that the binary system releases as the orbital separation between the core of the giant star and the companion decreases. This is likely to energize more the equatorial outflow than the polar outflow, either when the secondary star companion is outside the envelope, e.g., by ejecting mass through the second Lagrangian point (e.g., \citealt{Livioetal1979, MastrodemosMorris1999, Pejchaetal2016, Chenetal2017, Pejchaetal2017}) or during the common envelope evolution (e.g., \citealt{Iaconietal2017b, DeMarcoIzzard2017, Galavizetal2017, Iaconietal2017a, Chamandyetal2018, GarciaSeguraetal2018, MacLeodetal2018}, and references for older than two years papers therein). The second energy source is the gravitational energy that the mass that the secondary more compact star accretes from the envelope of the giant star releases. The secondary star can be inside the envelope of the giant star or outside the envelope. The most efficient process to carry the accretion energy to the outflow is to launch jets from an accretion disk. Such jets influence the outflow along and near the polar directions much more than they affect the equatorial outflow. In recent years the notion that in many cases jets shape the outflows from evolved stars, and in particular from AGB stars that evolve to become PNe, has benefited from progress in several directions. First is the realization that jets shape many PNe. Although the suggestions for jet-shaping of some PNe is old (e.g., \citealt{Morris1987, Soker1990AJ}), in the last two decades researchers have realized that jets shape a large fraction of non-spherical PNe (e.g., \citealt{SahaiTrauger1998, Boffinetal2012, HuarteEspinosaetal2012, Balicketal2013, Miszalskietal2013, Tocknelletal2014, Huangetal2016, Sahaietal2016, RechyGarciaetal2016, GarciaSeguraetal2016, Dopitaetal2018}). Then is the understanding that in many cases the jets operate in a feedback mechanism (see \citealt{Soker2016Rev} for a review). In the negative-feedback part the jets remove mass from the ambient medium which serves as the reservoir of accreted gas, hence reducing accretion rate that is followed by a decrease in jets' power (e.g., \citealt{Sokeretal2013, LopezCamaraetal2018}). The positive-feedback part comes from the removal of energy and gas from the very inner regions of the accretion disk, just near the accreting star. Consequently, this reduces the pressure in those regions and hence allows for a high accretion rate, even at super-Eddington rates (e.g., \citealt{Shiberetal2016, Chamandyetal2018}). The third direction of progress has been the finding that binary systems shape most, and probably all, PNe (e.g., \citealt{ Akrasetal2016, Alietal2016, Bondetal2016, Chenetal2016, Chiotellisetal2016, GarciaRojasetal2016, Hillwigetal2016a, Jonesetal2016, Madappattetal2016, Chenetal2017, DeMarcoIzzard2017, Hillwigetal2017,JonesBoffin2017, Sowickaetal2017, Alleretal2018, Barkeretal2018, Bujarrabaletal2018, Ilkiewiczetal2018, Jones2018, Miszalskietal2018b}, for a sample of papers from the last 3 years; for a different model see \citealt{GarciaSeguraetal2005}). In some cases there is a direct link between the presence of a binary central star and the presence of jets (e.g., \citealt{Boffinetal2012, Miszalskietal2013, Miszalskietal2018a}), and binary AGB systems and the presence of jets launched by the companion to the AGB star (e.g., \citealt{Thomasetal2013, Gorlovaetal2015, Bollenetal2017, VanWinckel2017}). The finding of binary systems in PNe is relevant to the present study because single AGB stars cannot launch jets for the lack of angular momentum, and the presence of a binary companion is crucial for launching jets in these progenitors of PNe. In light of the large number of observations that show that jets shape many PNe and other nebulae around evolved stars, we continue our study of the different morphological structures that jets can form. We stress here that by jets we refer to any bipolar outflow from the accretion disk, most likely in a binary system. This bipolar outflow can be two opposite narrow jets, a wide bipolar outflow from the accretion disk, it can be a continuous bipolar outflow or a chain of clumps. To all of these and more, we refer as jets. There are many simulations of jet-shaping of PNe and related nebulae (e.g., \citealt{LeeSahai2004, Dennisetal2009, Leeetal2009, HuarteEspinosaetal2012, Balicketal2013, Akashietal2015, Balicketal2017, Akashietal2018, Balicketal2018}), but here we concentrate on a particular morphological feature that we term columns crown. In the present paper we use the hydrodynamical core FLASH (section \ref{sec:numerical}) to study the formation of columns or filaments that protrude from lobes that jets inflate in the direction of the jets initial velocity. We describe the formation of the protruding columns, the columns crown, in section \ref{sec:results}. We are motivated by such thin columns that are observed in the PN Menzel~3 (Mz~3; PN~G331.7-01.0; the Ant nebula). In an earlier paper \citep{AkashiSoker2008a} we showed that a jet interacting with the AGB wind can form a lobe with a front lobe as observed in Mz~3. We now perform a simulation with a different set of initial conditions that lead to the formation of a delicate structure of the columns crown. The PN MZ~3 itself is a well studied PN (e.g., \citealt{Clyneetal2015, Alemanetal2018}) that we discuss more in section \ref{sec:summary} where we summarize our main findings. | \label{sec:summary} The purpose of this study is to enrich the variety of morphological features that jets can form when they interact with the slow outflow from evolved giant stars. The rich variety of morphological structures of PNe and other nebulae around evolved stars motivate us to conduct these hydrodynamical simulations. Our motivation to conduct this study is the beautifully complicated structure of the PN Mz~3, the Ant nebula. In Fig. \ref{fig:ant} we present this PN. \begin{figure} % \begin{center} \vskip -0.7 cm \includegraphics[trim= 1.5cm 16.0cm 2.0cm 3.0cm,clip=true,width=0.95\textwidth]{ColumnsCrownComparisonFig} \caption{Left panel: A false-color image of the PN Mz~3 by the Hubble Space Telescope. (Credit: NASA, ESA and The Hubble Heritage Team (STScI/AURA); Acknowledgment: Raghvendra Sahai and Bruce Balick.). Middle panel: An image of Mz~3 with marks from \cite{Clyneetal2015}. Right panel: Our numerical 3D density structure. } \label{fig:ant} \end{center} \end{figure} As with previous studies in this series, we set very simple initial conditions, i.e., a spherical slow ambient medium and one jet-launching episode. In the present study we set one dense shell within a slow wind from an AGB star, and injected two opposite jets in one episode that lasted $17 \yr$. In reality, the slow outflow is expected to be more complicated than we assume, and there can be several jet-launching episodes. For example, the jets do not need to encounter a closed shell, but rather it is sufficient that each jet catches up and interacts with a dense slow polar cap with a half opening angle of $\ga 25^\circ$. Such caps can be ejected by an earlier slow bipolar outflow. This will alleviate the non-detection of such spherical shells in PNe (e.g., \citealt{Sahaietal2011}). This type of interaction forms a bipolar nebula with two prominent morphological features on each of the two sides of the equatorial plane, as we mark on Fig. \ref{fig:Scematic}: (1) A bipolar lobe that is composed of two bubbles, and (2) a columns-crown. In Figs. \ref{fig:3D} and \ref{fig:dens_slice} - \ref{fig:vel_arrows} we presented the evolution of the interaction. After the jets cease to exit, the bipolar structure that the jets inflated continues to move forward. The interaction with the dense shell splits each lobe into two bubbles, one touching the center and one further out. The interaction at its several stages has many regions that are Rayleigh-Taylor unstable. In Fig. \ref{fig:RT} the instability maps show these unstable regions that form the columns crown and the filaments and clumps inside the outer bubble. Let us compare the bipolar structure that we obtained in our simulation with the structure of Mz~3. We summarize the comparison in Table \ref{table:Compare}. Although we terminate the simulation at an age of about 300 years, the outer parts of the nebula, including the columns crown, have very high Mach numbers, $\simeq 10$, and hence experience a ballistic motion that preserves the shape of the nebula. We therefore can apply our results to the older PN Mz~3. Each of the two lobes of Mz~3 is composed of two bubbles. However, these are not exactly the same as the two bubbles we obtained here. While in our simulation the outer bubble is larger than the inner one, in Mz~3 the outer one is smaller than the inner one. In a previous paper \citep{AkashiSoker2008a} we termed these outer small bubbles `front lobes'. In that paper we set different initial conditions, in particular there was no dense shell, and show how a jet can inflate a front lobe. Based on the earlier paper and the present study, we suggest one of the following two possibilities. (1) The exact setting of the dense ambient gas is more complicated, with some structure in between the setting we used in the two runs. (2) There were two jets-launching episodes in Mz~3, one episode that formed the two opposite crowns and one that formed the two front-lobes. We also note that a simple way to form two touching bubbles is to have two jets-launching episodes, one after the other. This requires no dense shell. \begin{table*}[t] \small \centering \caption{Observed and simulated properties of Mz~3.} \begin{tabular}{c c c c} % \hline Property & Observed & Simulated & Possible solutions for discrepancies\\ [0.5ex] \hline \hline General structure & Elongated &Reproduced & \\ & Bipolar nebula & & \\ \hline Columns-crown & Marked BL~2 &Reproduced & \\ &in fig. \ref{fig:ant}& & \\ \hline Opening of the & Very small (almost & $25^\circ$ to the & A more complicated initial \\ crown ¶llel to axis) & symmetry axis & ambient gas structure \\ \hline Velocity of & $\simeq 100 \km \s^{-1}$ & $\simeq 500 \km \s^{-1}$ & Slower jets\\ columns & & & \\ \hline Material near & Not observed. & Leftover from & Interaction closer to center \\ equatorial plane & & dense shell & followed by ballistic expansion\\ \hline Front lobe & clear on one & Not reproduced & A later jet launching episode\\ (Fig. \ref{fig:ant})& side & & as in \cite{AkashiSoker2008a} \\ \hline Two connected & Outer bubble is & Two bubbles, but & A somewhat different initial \\ bubbles (Fig. \ref{fig:ant}) & smaller & outer is larger & ambient medium structure \\ \hline Side rays & Marked BL~3 & Not reproduced & A jet-launching episode \\ &in fig. \ref{fig:ant}& & of wide jets\\ \hline Width of columns & Narrow & Wider than observed & Simulate with a higher \\ & & & resolution and a somewhat \\ & & & different ambient structure \\ \hline Origin of columns &Extended zone & A ring on the lobe. & Several short jets or a \\ &on the lobe. & & longer one jets-episode \\ \hline X-ray emission & Observed inside & Simulation has high & Late jets and a fast wind \\ & the lobes & Temperature gas; not& might improve the structural \\ & &exactly same structure& fitting \\ \hline \end{tabular} \label{table:Compare} \newline Comparison of the observed properties of Mz~3 with the results of the present simulation. \end{table*} Table \ref{table:Compare} list many properties of Mz~3 that we did not reproduce in the present simulation. There are two options to reconcile the discrepancies. The first is that our suggested jets-shell interaction is not the process that formed the columns crowns. The second one is that jets-shell interaction is the explanation to the columns-crowns, but there are several missing ingredients from our simulations. In the last column of Table \ref{table:Compare} we list these possible missing ingredients. The Table shows that we fail to reproduce most of the features of Mz~3. Future studies will have to examine other processes that can form the columns-crowns instead of jets, or else will reproduce with much more complicated simulations all properties of Mz~3, but leave jets as the explanation for the columns crown. The distance to Mz~3 is uncertain, with different studies listing values in the range of about $1-3 \kpc$ (see discussion by \citealt{Clyneetal2015}). At a distance of $1.5 \kpc$ the width of the main nebula we obtained in our simulation is about $20 ^{\prime \prime}$. This is compatible with the observed width of the main nebula (e.g., \citealt{Clyneetal2015}). \cite{SantanderGarciaetal2004} find the maximum velocity of the columns of the crown to be about $300 \km \s^{-1}$. The velocity of the columns that \cite{Clyneetal2015} observe (BL2 in their nomenclature), on the other hand, is only about $100 \km \s^{-1}$. We find that the velocity of the bullets in our simulation is about $500 \km \s^{-1}$, and the velocity of the column that trails each bullet is somewhat lower. This shows that we did not reproduce the structure of Mz~3 quantitatively, but only qualitatively, but not perfectly. To obtain a better match we will have to thoroughly study the parameter space. The two opposite columns crowns we have obtained here resemble the `crowns' in Mz~3, but they are not identical. Firstly, The `columns' in Mz~3 are much narrower than what we obtained here, and they better be termed filaments. Secondly, the `columns' in Mz~3 originate from an extended region on the lobes of Mz~3, while our columns originate from one ring on the outer bubble (Fig. \ref{fig:z=1e17}). We suggest to reconcile these differences as follows. The reason we obtain thick columns is because of the limited resolution of the grid. A much finer grid (beyond our capabilities) might form thiner columns. As for the origin of the filaments. We used a short jets-launching episode that lasted only 17 years. It is possible that a longer episode, or several short episodes, might lead to a more extended crown. There are other features in Mz~3 that we do not reproduce. One example is the columns at very large angles, marked as BL3 by \cite{Clyneetal2015}. Another is the exact structure of the X-ray emission within the main nebula \citep{Kastneretal2003}. Although we do obtain high temperature gas that emits X-ray, it is not clear it will survive for a long time, and the exact structure is not as that in Mz~3. Although we do not reproduce the exact X-ray structure, the X-ray structure does support shaping by jets \citep{Kastneretal2003} We attribute the more complicated observed nebular structure to two types of effects. The first is a much more complicated slow wind prior to the launching of jets, and the second is later jets-launching episodes. On top of these, the central star can blow a fast wind. Adding more features to the slow wind, and adding more jets-launching episodes make the parameter space much too large to follow. For that reason in our different studies we tend to study one type of interaction at a time. We summarize by stating that our main finding in the present study is that the type of Rayleigh-Taylor instability modes that develop in our simulation when jets and the lobes they inflate encounter circumstellar matter with large density gradients, larger than a wind with a constant mass loss rate, can account for the outward extending columns/filaments in Mz~3 and other nebulae. We thank Bruce Balick and an anonymous referee for enlightening and for useful comments. We acknowledge support from the Israel Science Foundation and a grant from the Asher Space Research Institute at the Technion. | 18 | 8 | 1808.00276 |
1808 | 1808.07660_arXiv.txt | Imaging of the dust continuum emitted from disks around nearby protostars reveals diverse substructure. In recent years, theoretical efforts have been intensified to investigate how far the intrinsic dynamics of protoplanetary disks (PPDs) can lead to such features. Turbulence in the realm of non-ideal magnetohydrodynamics (MHD) is one candidate for explaining the generation of zonal flows which can lead to local dust enhancements. Adopting a radially varying cylindrical disk model, and considering combinations of vertical and azimuthal initial net flux, we perform 3D non-ideal MHD simulations aimed at studying self-organization induced by the Hall effect in turbulent PPDs. To this end, new modules have been incorporated into the \NIR and \FTD MHD codes. We moreover include dust grains, treated in the fluid approximation, in order to study their evolution subject to the emerging zonal flows. In the regime of a dominant Hall effect, we robustly obtain large-scale organized concentrations in the vertical magnetic field that remain stable for hundreds of orbits. For disks with vertical initial net flux alone, we confirm the presence of zonal flows and vortices that introduce regions of super-Keplerian gas flow. Including a moderately strong net-azimuthal magnetic flux can significantly alter the dynamics, partially preventing the self-organization of zonal flows. For plasma beta-parameters smaller than 50, large-scale, near-axisymmetric structures develop in the vertical magnetic flux. In all cases, we demonstrate that the emerging features are capable of accumulating dust grains for a range of Stokes numbers. | \label{sec:intro} Young stellar objects are found to harbor protoplanetary accretion disks (PPDs) composed of partially ionized gas and dust grains inherited from the interstellar medium. Renewed theoretical efforts have aimed at establishing a sound picture of the structure and evolution of these objects \citep[see][for a recent review]{Turner2014PPVI}. These efforts are of paramount importance both for providing a framework in which theories of planet formation can be developed, and for aiding the interpretation of observations in the infrared and (sub-) millimeter wavebands \citep[e.g.,][]{2011ARA&A..49...67W}. Observations appear to indicate that dust is settled-out to the disk midplane with a scale-height of $1\au$ at an orbital location of $100\au$ \citep{Pinte2016}. In addition, they suggest that grains grow faster in the inner regions of the disk or, alternatively, that larger grains drift inward more efficiently. In view of the growth of the embedded dust --- first to the mm-sizes detected in thermal emission in nearby PPDs and witnessed as chondrules in solar system meteorites, and then further to cm-sized pebbles and ultimately planetesimals --- it is important to consider the coupled aerodynamic evolution of the dust and the gas. In featureless disks, intermediate-size dust grains are prone to rapid depletion due to their swift radial drift toward the star \citep{Whipple,Weidenschilling}. Due to the drag force, gas pressure variations can slow down or stall the radial drift, creating regions of increased dust density, sometimes referred to as ``dust traps.'' Such pressure gradients can arise near the orbital location of (already formed) planets that are perturbing the gas disk, at opacity transitions, ice lines, dead-zone boundaries, or spontaneously via nonlinear feedback, as well as via self-organization of the flow into zonal flows or vortices \citep[see][for recent reviews]{Johansen2014, Testi2014PPVI}. Outflows in the form of collimated protostellar jets and wide-angle disk winds \citep{Arce2007PPV,Bjerkeli2016} point toward the presence of dynamically relevant magnetic fields. In particular, in the context of magnetized, turbulent disks resulting from the magnetorotational instability \citep[MRI;][]{Balbus1991}, pressure gradients can arise via the emergence of zonal flows \citep{Johansen2009}, or large-scale vortical flows \citep{Fromang2005,Johansen2005}. Shearing box simulations of MRI in the ideal-MHD setting carried out by \citet{Johansen2009} showed that axisymmetric surface density fluctuations can grow and persist for many orbits. Moreover, the resulting ``pressure bumps'' are typically found to be in geostrophic balance and are thus accompanied by sub- / super-Keplerian rotation. By means of 3D global simulations including Ohmic diffusion, \citet{Dzyurkevich2010} concluded that the excavation of gas from the active region during the linear growth and after the saturation of the MRI leads to the creation of a steady local radial gas pressure maximum near the dead zone edge, as well as to the formation of dense rings within the MRI-active region. Beyond their very inner reaches (i.e., at radial distances from the star larger than a few tenths of an AU), PPDs are too cold for the gas to be thermally ionized \citep{DeschTurner2015}. At the same time, from a few AU to a few tens of AU, the interior of the disk, near its midplane, is efficiently shielded from exterior sources of ionizing radiation \citep{TurnerDrake2009}, leaving the bulk of the gas in a state of partial ionization. Thus, it is necessary to study the evolution of magnetic fields in the framework of \emph{non-ideal} MHD, where dissipative processes due to collisions are taken into account. The central importance of the non-ideal effects has long been recognized \citep[see, e.g.][]{Wardle1999,Sano2002,Salmeron2003}, but it is only recently that global simulations including all non-ideal effects --- i.e., Ohmic resistivity, ambipolar diffusion (AD) and the Hall term --- have been performed in relevant parameter regimes. This has led to a modified picture of how PPDs accrete \citep{Bai2013, Gressel2015,Bai2017,Bethune2017}, a picture that deviates significantly from the traditionally envisioned dead zone \citep{Gammie1996}. Detailed calculations of chemical abundances and ionization fractions using chemical reaction networks have provided an understanding of the relative importance of the three non-ideal effects \citep{Salmeron2008,Bai2011a}. Ohmic resistivity becomes important in regions of high density and weak magnetic fields. Those are typically expected to be located in the midplane region of the PPD, between approximately 1 and 5 AU. The Hall effect is estimated to dominate close to the midplane, between 1 and 10 AU, in what is sometimes referred to as the planet-forming regions of PPDs. AD, in turn, is anticipated to dominate in low-density regions with comparatively strong magnetic fields -- conditions applicable to the surface layers of the inner reaches of the disk, as well as the bulk of the outer disk (i.e., $\simgt 30\au$). The role of AD in the formation of zonal flows and its potential impact on halting the inward drift of dust particles were first studied by \citet{Simon2014} using the local shearing-box approximation. Subsequently, using 3D unstratified global simulations, \citet{Zhu2014} found that turbulence is significantly suppressed around a planet-carved gap, and that a large vortex can form at the edge of the gap. The vortex can efficiently trap dust particles that span three orders of magnitude in size within a timescale of about 100 orbits at the planet's location \citep{Zhu2015}. Shearing box simulations including AD have been employed by \citet{Bai2014dust}, who found that magnetic flux can be concentrated into thin axisymmetric shells, whose typical extent is less than half a pressure scale height. Non-ideal global 3D MHD stratified simulations of the dead-zone outer edge \citep{Flock2015} have been analyzed in terms of potential observational signatures by \citet{Ruge2016}. More recently, \citet{Suriano2017} have discussed the problem of dust trapping from the perspective of a magnetized wind in an AD-dominated disk. In the presence of the Hall effect, the segregation of the turbulent flow into zonal bands appears to be even more pronounced. For a local, vertically unstratified model, \citet{Kunz2013a} found that in Hall-dominated magnetorotational turbulence, zonal flows with radial concentrations of vertical magnetic flux develop. The zonal flows are characterized by strong field amplitudes and are driven by a coherent Maxwell stress acting in concert with conservation of canonical vorticity --- a generalization of the flow vorticity motivated by the additional velocity appearing in the Hall-MHD induction equation. Recent work by \citet{Bethune2016a} has confirmed the self-organization in a global cylindrical model, but still radially and vertically unstratified. The authors reported the generation of large-scale vortices and zonal flows, suggesting the possibility of dust trapping in the produced features. Within the shearing-box framework, \citet{Lesur2014} showed that the Hall effect can induce a strong azimuthal field when vertical stratification is considered and zonal flows related to the local confinement of vertical magnetic flux can be inhibited. \citet{Bai2015} also found zonal flows of vertical magnetic field in unstratified shearing box models; meanwhile stratified simulations show thin zonal flows that are supposedly not produced by the Hall effect. The goal of our work is to assess the existence of dust traps in the form of zonal flows and vortices induced by Hall-MHD turbulence in protoplanetary disks. We focus our current modeling efforts to the case of a vertically unstratified cylindrical disk considering only Ohmic dissipation and the Hall effect. Such a simplification can be safely applied to the midplane of a typical PPD between $1 \textrm{ and } 5\au$, where the Hall effect likely dominates over Ohmic dissipation, and where AD is thought to be negligible. We first aim at confirming the existing work and then extend the scope of the models to the dynamics of a relatively strong toroidal field ($B_{\phi0} \sim 10 B_{z0}$). We furthermore include dust fluids to rigorously confirm the ability of Hall-modified MRI turbulence to create self-organized features that are able to trap dust. This paper is organized in the following manner. In | , and document our code implementation, as well as present different benchmarks in Appendices~\ref{sec:apA} and \ref{sec:apB}, respectively. \label{sec:conclusion} We have demonstrated that our implementations of Hall-MHD in \NIR and \FTD are suitable for studying problems in the context of protoplanetary disks subject to the non-ideal MHD effects. The inclusion of an artificial resistivity is only necessary when the Hall effect strongly dominates in the induction equation by a factor more than $10^2$, with respect to the ideal-MHD induction term. In all the models explored in this strong Hall regime, around 5\% of the active domain has an artificial magnetic Reynolds number $1<\Rma<5$, so we conclude that the dissipation likely has a minimal impact on the dynamics and that it overall does not affect the results that we have obtained. A central new result of our investigation is that the self-organization is sustained in a disk with radial stratification. This is also true if the azimuthal flux is comparable to the vertical flux for a plasma-$\beta$~parameter in the range between $\beta_{z} = 10^4 - 10^3$. Simulations where the azimuthal domain is reduced to $L_{\phi} = \pi/2$ show axisymmetric zonal flows and vortices, meanwhile we are only able to recognize large-scale vortical features if the domain is $L_{\phi}\eq2\pi$. We confirm the finding that the zonal flows (and vortices) are confined between regions of strong Maxwell stress --- in agreement with the mechanism described by \citet{Kunz2013a}. Because of the low flow compressibility and the inclusion of small Ohmic dissipation, the vortices and zonal flows are anti-correlated with low vorticity regions, which is expected when the magneto-vorticity is locally conserved. When the initial model has a strong azimuthal net flux, i.e., $\beta_{\phi} \simeq 10^{-2}\tms \beta_{z}$, azimuthally large-scale concentrations of vertical magnetic flux are still possible to obtain. However, the picture of field confinement between regions of enhanced Maxwell stress is no longer clearly identified in these simulations. The azimuthal component of the magnetic field strongly dominates the dynamics and the self-organized zonal flows are harder to recover, which is in agreement with the previous results by \citet{Lesur2014} and \citet{Bethune2017}. The inclusion of a net-azimuthal flux does not alter the evolution of the zonal flows when $\beta_{\phi}\sim \beta_z$ \citep[also cf.][]{Bethune2016a}. The spectral energy distribution increases as we move towards shorter azimuthal wavenumber, $m \simeq 4$. When $\beta_{\phi}=5\ee{3}$, the distribution shows a maximum followed by a flat profile at the lower azimuthal modes. Increasing $\beta_{\phi}$ leads to a spectrum that has a peak at intermediate azimuthal wavenumbers. These energy distributions have been obtained using a disk with a reduced azimuthal extent of $\pi/2$. The fact that the maxima are located close to the lowest azimuthal wavenumbers implies that a full $2\pi$ domain is advisable for studying the global evolution of the zonal flows. By including the evolution of pressureless dust fluids in the ensuing Hall-MHD turbulence, we demonstrate that quasi-axisymmetric dust enhancements can be obtained for the range of Stokes numbers explored -- even in the absence of prominent flow features, such as the result of Hall-effect induced self-organization. There appears to be a difference in character regarding the precise nature of the ensuing dust traps, however. In models \emph{without} azimuthal net flux, dust enhancements are typically located at the position of vortices themselves, which agrees well with the notion of vortices being able to trap dust \citep{1997Icar..128..213K}. In contrast to this, in models with significant azimuthal magnetic flux ($\beta_\phi\simeq 50$), the dust accumulations appear to coincide with local concentrations of the vertical magnetic flux. For particles with Stokes number $\St\eq1$, peaks in the dust concentration with $\rho_{d}/\rho_{d0} \simeq 20$ are reached during the integration time of a few hundred inner orbits that we have adopted here, implying that the feedback onto the gas might have to be considered when pursuing longer integration times. The drift time scale of the $\St=0.1$ and $\St=0.01$ particles can be estimated to become between one and two orders of magnitude higher than that for $\St=1$. As a result, lower dust enhancement factors were obtained for these two dust species, which we, at least partially, attribute to the insufficient evolution time covered by our current simulations. In conclusion, we have demonstrated the ability of our numerical implementations to deal with the MHD turbulence created by the Hall-MRI in the regime of a dominant Hall effect, and were able to study its consequences for the local accumulation and segregation of intermediate-Stokes number dust grains. The presented numerical tools will enable us to further explore the phenomenon of self-organization in the context of Hall-MHD turbulence, as well as its potential relevance in explaining the axisymmetric and non-axisymmetric substructure observed in nearby protoplanetary disks in the (sub-)mm waveband. \software{ NIRVANA-III \citep{2004JCoPh.196..393Z,2011JCoPh.230.1035Z,2016A&A...586A..82Z}, FARGO3D \citep{Benitez-Llambay2016}, IPython \citep{numpy}, Matplotlib \citep{matplt} } | 18 | 8 | 1808.07660 |
1808 | 1808.05386_arXiv.txt | We study the radial metallicity gradient $\Delta[M/H]/\Delta R_g$ as a function of [Mg/Fe] and $|Z|$ with the help of a guiding radius based on the Apache Point Observatory Galactic Evolution Experiment and Gaia and then analyze the radial migration effect on the radial metallicity gradient and metallicity-rotation gradient between the Galactic thin and thick disks. The derived trend of gradient $\Delta[M/H]/\Delta R_g$ versus [Mg/Fe] shows a transition at [Mg/Fe] $\sim 0.18$ dex, below which the gradient is negative and varies a little as [Mg/Fe] increases; however, it changes sharply in [Mg/Fe] ranges of 0.16-0.18, above which the gradient increases linearly with increasing [Mg/Fe], being a positive value at [Mg/Fe]$\gtrsim 0.22$ dex. These positive gradients in the high-[Mg/Fe] populations are found at $|Z| < 0.8$ kpc, and there are nearly no gradients toward higher $|Z|$. By comparing the metallicity distributions, the radial metallicity gradients $\Delta[M/H]/\Delta R$ and the metallicity-rotation gradients between the total sample and $|R-R_g|<2$ kpc subsample (or $|R-R_g|>2$ kpc subsample), we find that, for the thick disk, blurring flattens the gradient $\Delta[M/H]/\Delta R$ and favors metal-poor high-eccentricity stars. These stars are responsible for the measured positive metallicity-rotation gradient of the thick disk. | The Milky Way disk has been suggested to host a thick disk, in addition to its thin disk, based on tje results of geometric decomposition works \citep[e.g.][]{Gilmore-Reid1983,Juric2008,Jia2014}. Generally, the thick disk is thought to be old as the thin disk \citep[e.g.][]{bensby05,yoachim08}. Other age-related properties, such as $\alpha$-element abundances \citep[e.g.][]{fuhrmann1998, bensby03,bensby05, reddy06, haywood08,Adibekyan2011} and kinematics \citep{Chiba2000,Gilmore2002,Yoachim05}, are also found to be different from the thin disk. Although the formation mechanism of the thick disk is still debated, four kinds of scenarios are widely known so far: heating \citep[e.g.][]{Quinn1993}, accretion \citep[e.g.][]{Abadi2003}, gas-rich mergers \citep[e.g.][]{Brook2004,Bournaud2009}, and radial migration \citep[e.g.][]{Schonrich-Binney2009}. In particular, the last one has attracted much attention because it could be triggered by the well-known non-axisymmetric structures of the Milky Way disk (such as spiral arms and bars). On the other hand, radial migration would alter the disk structure and chemical composition through churning (stars changing angular momentum) and blurring \citep[stars conserving their angular momentum,][]{Schonrich-Binney2009}, which makes the interpretation of the observational results complicated. In this respect, the study on the radial metallicity gradient, which records the enrichment history of stellar abundances at a given radial location, provides a good way to disentangle the radial migration effect on the chemical evolution of the Galactic disk. The radial metallicity gradient has been well studied in previous works based on a variety of stellar tracers. \citet{Tas2016} gave a detailed summary in their Table 1, where the thin disk has a negative gradient and the thick disk shows a flat or even positive gradient. Moreover, the overall disk gradient was found to gradually flattens toward high $|Z|$, which is proposed to be driven mainly by radial migration \citep{Schonrich2017}. Interestingly, the metallicity-rotation gradient is also found to be negative for the thin disk but positive for the thick disk \citep[e.g.][]{Adibekyan2013, Allende-Prieto2016}, which is also linked with the radial metallicity gradient. It has been suggested that the negative (positive) metallicity-rotation gradients is driven by the negative (positive) radial metallicity gradient for thin (thick) disk stars due to blurring effect \citep{Vera-Ciro2014,Allende-Prieto2016,Schonrich2017}. In this case, a thick disk star born in the inner disk with epicyclic motion would tend to be metal poor due to the positive radial metallicity gradient. Meanwhile, its current location in the solar neighborhood is near the orbit's apo-center and thus tends to rotate slower. This naturally results in a positive metallicity-rotation gradient for thick disk stars. Supporting this scenario, \citet{Toyouchi2014} reported that the measured radial metallicity gradients for $\alpha$-enhanced (thick disk) populations are positive. Note that \citet{Toyouchi2014} adopted the guiding radius $R_g$, instead of current Galactocentric distance $R$, since the gradient, if measured with $R$, likely suffers from the radial migration effect and the usage of the guiding radius $R_g$ can diminish the blurring effect. However, the elemental abundances in \citet{Toyouchi2014} are based on low-resolution spectra of the Sloan Digital Sky Survey (SDSS), and distances are based on photometric methods. These can be greatly improved using recently released high-quality data. Investigating the radial metallicity gradient by taking into account the guiding radius and the metallicity-rotation gradients based on accurate elemental abundances and radial velocity from high-resolution spectra provided by the Apache Point Observatory Galactic Evolution Experiment \citep[APOGEE,][]{Majewski2017} survey is of much interest. Additionally, the wealth astrometric information from Gaia \citep{Gaia2016} enables us to perform a statistical analysis on the radial metallicity gradient with the aid of the guiding radius. Specifically, we aim to derive the radial metallicity gradient as a function of [Mg/Fe] and $|Z|$ and analyze the effect of radial migration on the radial metallicity gradient and metallicity-rotation gradient between the Galactic thin and thick disks. | In this work, we investigate the radial metallicity gradient between the thin and thick disks, with the help of the guiding radius, and analyze the effect of radial migration on the derived radial metallicity gradient and the metallicity-rotation gradient based on a large and high precision data from Gaia and APOGEE. To infer the Galactic disk formation history, we derive the radial metallicity gradient $\Delta[M/H]/\Delta R_g$ as a function of [Mg/Fe] and $|Z|$. We find that the radial metallicity gradient of low-[Mg/Fe] population ([Mg/Fe]$<$0.16 dex) is negative and has small variations with [Mg/Fe], and this gradient changes sharply until [Mg/Fe] reaches 0.18 dex, after which the gradient rises linearly and becomes positive at [Mg/Fe]$\ge 0.22$. The positive gradients in the high-[Mg/Fe] population are found to limit at $|Z|<0.8$ kpc only, and there is nearly no gradient toward higher $|Z|$. The above results imply that the Galactic disk formed in an inside-out and upside-down fashion. Moreover, the different trends of $\Delta[M/H]/\Delta R_g$ versus [Mg/Fe] between the thick disk and thin disk, which are characterized by the high-[Mg/Fe] population ([Mg/Fe]$>$0.18 dex) and the low-[Mg/Fe] population ([Mg/Fe]$<$0.18 dex), respectively, indicate that the transition between them occurs at [Mg/Fe]$\sim$0.18 dex. With a more specific division of the thin and thick disks in the chemical plane, we still find that the trends of $\Delta[M/H]/\Delta R_g$ versus [Mg/Fe] are rather different between the thin and thick disks, which indicates that the thin disk and thick disk are two distinct star populations. In the meantime, the thick disk clearly presented a flat or positive gradients depending on the [Mg/Fe] ratio. In order to study the radial migration effect, we adopt a quantity, $|R-R_g|$, to reflect the blurring effect since blurring does not change on $R_g$ but does on $R$. A subsample with $|R-R_g|<2$ kpc is thought to not suffer from the blurring effect, and the remaining subsample with $|R-R_g|>2$ kpc is considered to contain migration stars. By comparing the metallicity distributions, the radial metallicity gradients $\Delta[M/H]/\Delta R$, and the metallicity-rotation gradients between the total sample and the alternative two above subsamples, we find that blurring process flattens the gradient $\Delta[M/H]/\Delta R$ and favors metal-poor high-eccentricity stars for the thick disk. These stars are responsible for the measured positive metallicity-rotation gradient in the thick disk. | 18 | 8 | 1808.05386 |
1808 | 1808.02045_arXiv.txt | The ionization mechanism of the low-ionization gas in quiescent red sequence galaxies has been a long-standing puzzle. Direct temperature measurements would put strong constraints on this issue. We carefully selected a sample of quiescent red sequence galaxies from SDSS. We bin them into three bins with different \nii/\hal\ and \nii/\oii\ ratios, and we measure the temperature-sensitive \oiiitw, \niitw, \siitw, and \oiitw\ lines in the stacked spectra. { The \sii\ doublet ratios indicate the line-emitting gas is in the low density regime ($\sim10-200$ cm$^{-3}$).} We found the temperatures in the S$^+$ zones to be around 8000K, the temperatures in the O$^+$ zone to be around $1.1-1.5\times10^4$K, and the temperatures in the N$^+$ zones to be around $1-1.4\times10^4$K. { The \oiiitw\ line is not robustly detected.} We found that the extinction corrections derived from Balmer decrements would yield unphysical relationships between the temperatures of the S$^+$ zones and O$^+$ zones, indicating that the extinction is significantly overestimated by the measured Balmer decrements. We compared these line ratios with model predictions for three ionization mechanisms: photoionization by hot evolved stars, shocks, and turbulent mixing layers. { For the photoionization and shock models, the hot temperatures inferred from \sii\ and \nii\ coronal-to-strong line ratios require metallicities to be significantly subsolar. However, the \nii/\oii\ line ratios require them to be supersolar. None of the models could simultaneously explain all of the observed line ratios, neither could their combinations do.} | % Contrary to the conventional wisdom, recent studies have shown that elliptical and lenticular galaxies also have interesting interstellar medium. They not only have X-ray-emitting hot gas, many of them also have substantial amount of warm and cold gas, and dust \citep[e.g.][]{Phillips86,Kim89,Buson93,Goudfrooij94,Macchetto96,Zeilinger96,Lauer05,Sarzi06,DavisT11, Singh13, Gomes16}. The evidence for the warm gas in early-type galaxies come from optical emission line measurements. Unlike late-type galaxies, in most cases, the emission lines are not due to star formation. Most early-type galaxies with line emission show a line ratio pattern similar to Low-ionization Nuclear Emission-line Regions (LINER, \citealt{Heckman80}). This often leads to them being classified as a type of active galactic nuclei. However, many recent studies have shown that these line emission is not only spatially-extended, but have a spatial gradient in line ratio that are consistent with being photoionized by sources that are spatially distributed like the stars, rather than a central AGN \citep{Sarzi10, YanB12, Singh13, Belfiore15, Gomes16}. Therefore, the name LINER is no longer appropriate. In some recent papers, these phenomena have been referred to as LINER-like. \cite{Belfiore15} recently has renamed them as Low-Ionization Emission Regions or LIERs for short. LINER-like and LIER are both referring to the same phenomenon, namely the spatially-extended optical emission lines observed in a large fraction of quiescent red galaxies where the line ratios do not match those of star-forming regions. The exact ionization mechanism certainly could be different in individual galaxies. But for the majority of early-type galaxies where it is dominated by extended emission, having uniform line ratio and similar equivalent width (EW), there is likely a common mechanism shared by most of them. Although the AGN photoionization model has been ruled out for the majority of low-ionization emission-line galaxy, there are several mechanisms that are spatially distributed, such as photoionization by hot evolved stars, collisional ionization by shocks, and turbulent mixing layers. Photoionization by hot evolved stars is a potential explanation for these warm ionized gas. It was originally proposed by \cite{Binette94}, and further developed by \cite{Stasinska08} and \cite{CidFernandes11}. Several observational evidences favor this explanation. The strongest evidences are two. One is the reasonably tight correlation between emission-line surface brightness and the stellar surface density shown by \cite{Sarzi10}. The other is that the ionization parameter gradient measured statistically by \cite{YanB12} matches the prediction of a distributed ionizing source following the stellar density profile. The latter work also found the luminosity-dependence of the ionization parameter gradient also match the theoretical prediction. These seem to be very strong evidence. However, there are also serious problems associated with this explanation. % First, the ionization photon budget does not work out. Current stellar evolution theory suggests the number of ionizing photons from an evolved stellar population is roughly $10^{41}$ per second per solar mass, and it changes little with age after 1Gyr. In order to power all the line emission, the gas has to absorb 100\% of the ionizing photons from post-AGB stars. This would require the gas to surround each individual stars. However, in many galaxies, the gas are observed to have different kinematics from the stars \citep{Sarzi06, DavisT11, Gomes16}, suggesting an external origin. If the neutral gas originated externally, we do not expect them to provide complete covering fraction. Thus, we cannot explain the level of emission given the current prediction of the stellar evolution theory. Second, if the gas originated externally and is randomly positioned relative to the hot evolved stars, we will also have trouble to produce the kind of ionization parameter that are seen. For detailed calculations, see \cite{YanB12}. Shocks are another popular mechanism for the ionization. Given the amount of stellar mass loss and the large stellar velocity dispersion in these early-type galaxies, cloud-cloud collisions will produce shocks. But it is unclear whether it is the dominant emission-line luminosity contributor. In an early-type galaxy observed by the SDSS-IV MaNGA (Mappping Nearby Galaxies at Apache Point Observatory) survey \citep{Bundy15,Yan16b}, \cite{Cheung16} found narrow bi-symmetric features in its \hal\ equivalent width map. The proposed theory is that winds powered by a central radio AGN produces shocks in the direction of the winds, enhancing the line emission in those directions. % Turbulent mixing layers may also produce the observed emission lines \citep{Slavin93}. Shear flows at the boundaries of hot and cold gas could produce gas at intermediate temperatures and give similar line ratios as observed in LIER and diffuse ionized gas in star-forming galaxies. One reliable way to distinguish these different ionization mechanisms is the temperature of the gas. Photoionized gas usually has a temperature around $10^4$K, while shocked gas and turbulent mixing has higher temperatures, approaching $10^5$K. If we can measure the temperature, we can tell which ionization mechanism is at work. There are a few temperature-sensitive line ratios in an optical spectrum, such as \oiiibw/\oiiitw, \niibw/\niitw. The weaker line in these line ratios are usually too weak to be detected in an SDSS-quality spectrum. In this paper, we measure these temperature-sensitive lines in carefully stacked spectra of quiescent red galaxies. This allows us to measure the gas temperature and put strong constraints on the ionization mechanisms. The paper is organized as the following. We describe the data in Section 2. The methods of sample selection, line measurements and zero-point corrections, selection of subsamples, stacking, stellar continuum subtraction, and weak line measurements in the stacked spectra are described in Section 3. We derive the extinction and gas temperatures in Section 4. We compare the results with models in Section 5 and conclude in Section 6. Throughout the paper, we assume a flat $\Lambda$CDM cosmology with $\Omega_m=0.3$ and a Hubble constant $H_0=100h$~km~s$^{-1}$ Mpc$^{-1}$. The magnitudes used are all in the AB system. | Using the spectra stacking technique, we have obtained high signal-to-noise emission-line spectra of quiescent red sequence galaxies from the SDSS. After careful continuum subtraction using a control sample without emission lines, we have detected multiple temperature-sensitive coronal emission lines or provided meaningful upper limits. { From reliable measurements of} \siitw, \oiitw, and \niitw\ lines, we { inferred that the emission from these ions was from regions} with a temperature of around $10^4$K. The measured \siitw/\siiw\ and \oiitw/\oiiw\ line ratios provide us with interesting constraints on the levels of extinction which are much lower than the values estimated from Balmer decrements. If the extinction values estimated from Balmer decrements { turn out to be the correct ones}, they would lead to unphysical temperature relationship between the S$^+$ and O$^+$ emission zones. The unresolved discrepancy between the two extinction estimates is one of the puzzles we uncovered in this paper. The \siibw/\siiaw\ ratio indicates { an electron density of} the line-emitting S$^+$ gas { that is smaller than or of the order of} 100 cm$^{-3}$. This result was used to set the physical conditions for our photoionization models { or to limit shock models either to those characterized} with a very low preshock density $n\le 1$cm$^{-3}$ and/or those with a strong magnetic field, since only such shock models can reproduce the observed \sii\ doublet ratios. We compared the temperature-sensitive line ratios with model predictions for three ionization mechanisms: photoionization by hot evolved stars, radiative shocks, and turbulent mixing layers. We found that neither the photoionization models nor the shock models can simultaneously explain all the line ratios we observe. Photoionization models with solar and supersolar metallicities can account for the strong line ratios. { On the other hand,} all of the temperature-sensitive line ratios indicate high temperatures, { which would imply} subsolar metallicities. The marginally-detected \oiiitw\ line in the high-\nii/\hal\ sample, if it is real, would indicate too high a temperature to be fit at all by photoionization models. Shock models, which have one more degree of freedom than photoionization models, similarly cannot explain all the line ratios. Among models that { reproduce the observed} low S$^+$ density, those with solar metallicity, n=10 cm$^{-3}$, and strong magnetic field can match the observed temperatures in O$^+$, S$^+$, and O$^{++}$ zones, but fail to match the high temperature of the N$^+$ zone. Lower pre-shock density ($n\le 1 {\rm cm}^{-3}$) shock models require significantly subsolar metallicities to match the temperature of the N$^+$ and S$^{+}$ zones, but would require supersolar metallicity in order to match the \nii/\oii\ ratios. The main discrepancy is between the observed high \nii/\oii\ ratios and the relatively hot temperatures we derived from coronal lines of singly ionized ions. Neither photoionization nor shock models could { reproduce both temperatures observed}. Given that both photoionization and shock models are failing to account for all observed line ratios for similar reasons, the combination of the two processes would face { similar difficulties.} { We could not do full justice to turbulent mixing models given that model predictions were only available for a few lines. } Our work illustrates the powerful constraints provided by temperature-sensitive line ratios. They reveal significant discrepancies between data and models. Although the interpretation { of our data might be questioned on the ground of the inevitable level of uncertainty associated with} the stacking process whereby different galaxies are averaged together, it encourages us to push deeper with { deeper observations that would aim at detecting the reported lines} in individual galaxies or, { if a significantly improved S/N was achieved}, would allow us to reduce the number of galaxies required for detecting them. Our work also highlights the need of more detailed and consistent modeling that would provide us with more stringent comparisons with observations when better data become available. | 18 | 8 | 1808.02045 |
1808 | 1808.00871_arXiv.txt | The fully analytical solution for isothermal Bondi accretion on a black hole (MBH) at the center of two-component Jaffe (1983) galaxy models is presented. In a previous work we provided the analytical expressions for the critical accretion parameter and the radial profile of the Mach number in the case of accretion on a MBH at the center of a spherically symmetric one-component Jaffe galaxy model. Here we apply this solution to galaxy models where both the stellar and total mass density distributions are described by the Jaffe profile, with different scale-lengths and masses, and to which a central MBH is added. For such galaxy models all the relevant stellar dynamical properties can also be derived analytically (Ciotti \& Ziaee Lorzad 2018). In these new models the hydrodynamical and stellar dynamical properties are linked by imposing that the gas temperature is proportional to the virial temperature of the galaxy stellar component. The formulae that are provided allow to evaluate all flow properties, and are then useful for estimates of the scale-radius and the mass flow rate when modeling accretion on massive black holes at the center of galaxies. As an application, we quantify the departure from the true mass accretion rate of estimates obtained using the gas properties at various distances from the MBH, under the hypothesis of classical Bondi accretion. | Observational and numerical investigations of accretion on massive black holes (hereafter MBH) at the center of galaxies often lack the resolution to follow gas transport down to the parsec scale. In these cases, the {\it classical} Bondi (1952) solution for spherically-symmetric, steady accretion of a spatially infinite gas distribution onto a central point mass is then commonly adopted; this is the standard reference for estimates of the accretion radius (i.e., the sonic radius), and the mass accretion rate (see, e.g., Rafferty et al. 2006, Sijacki et al. 2007; Di Matteo et al. 2008; Gallo et al. 2010; Pellegrini 2010; Barai et al. 2011; Bu et al. 2013; Cao 2016; Volonteri et al. 2015; Choi et al. 2017; Park et al. 2017; Beckmann et al. 2018; Ram\'irez-Velasquez et al. 2018; Barai et al. 2018). Even though highly idealized, during phases of moderate accretion (in the ``maintainance'' mode), indeed, the problem can be considered almost steady, and Bondi accretion could provide a reliable approximation of the real situation (e.g., Barai et al. 2012, Ciotti \& Ostriker 2012). However, leaving aside the validity of the fundamental assumptions of spherical symmetry, stationarity, and optical thinness, two major problems affect the direct application of the classical Bondi solution, namely the facts that 1) the boundary values of density and temperature of the accreting gas should be evaluated at infinity, and 2) in a galaxy, the gas experiences the gravitational effects of the galaxy itself (stars plus dark matter) all the way down to the central MBH, and the MBH gravity becomes dominant only in the very central regions, inside the so-called ``sphere-of-influence''. The solution commonly adopted in numerical and observational applications to alleviate these problems is to use values of the gas density and temperature ``sufficiently near'' the MBH, thus assuming that the galaxy effects are negligible. Of course, as the density and temperature of the accreting gas change along the pathlines, also the predictions of classical Bondi accretion will change, when based on the density and temperature measured at a finite distance from the MBH. It is therefore of great interest to be able to quantify the systematic effects on the accretion radius and the mass accretion rate obtained from the classical Bondi solution, due to measurements taken at finite distance from the MBH, and under the effects of the galaxy potential well. A first step towards a quantititative analysis of this problem was carried out in Korol et al. (2016, hereafter KCP16) where the Bondi problem was generalized to the case of mass accretion at the center of galaxies, including also the effect of electron scattering on the accreting gas. KCP16 then calculated the deviation from the true values of the estimates of the Bondi radius and of the mass accretion rate, due to adopting as boundary values for the density and temperature those at a finite distance from the MBH, and assuming the validity of the classical Bondi accretion solution. In the specific case of Hernquist (1990) galaxies, KCP16 obtained the analytical expression of the critical accretion parameter, as a function of the galaxy properties and of the gas polytropic index $\gamma$. However, even for this quite exceptional case, the radial profiles of the hydrodynamical variables remained to be determined numerically. Following KCP16, Ciotti \& Pellegrini (2017, hereafter CP17) showed that the whole accretion solution can be given in an analytical way (in terms of the Lambert-Euler $W$-function) for the {\it isothermal} accretion in Jaffe (1983) and Hernquist galaxy models with central MBHs. This meant that for these models not only it is possible to express analytically the critical accretion parameter, but also that the whole radial profile of the Mach number (and then of all the hydrodynamical functions) can be explicitely written. At the best of our knowledge, CP17 provided the first fully analytical solution of the accretion problem on a MBH at the center of a galaxy. The galaxy models used in KCP16 and CP17, i.e., the Hernquist and Jaffe models, are not only relevant because for them it is possible to solve the accretion problem, but also because of their numerous applications in Stellar Dynamics. In fact, these models belong to the family of the so-called $\gamma$-models (Dehnen 1993, Tremaine et al. 1994) and are known to reproduce very well the radial trend of the stellar density distribution of real elliptical galaxies; at the same time, their simplicity allows for analytical studies of one and two-component galaxy models (e.g., Carollo et al. 1995, Ciotti et al. 1996, Ciotti 1999). In particular, Ciotti \& Ziaee Lorzad (2018, hereafter CZ18), expanding a previous study by Ciotti et al. (2009), presented spherically symmetric two-component galaxy models (hereafter JJ models), where the {\it stellar} and {\it total} mass density distributions are both described by the Jaffe profile, with different scale-lengths and masses, and a MBH is added at the center. The orbital structure of the stellar component is described by the Osipkov-Merritt anisotropy (Merritt 1985). Moreover, the dark matter halo (resulting from the difference between the total and the stellar distributions) can reproduce remarkably well the Navarro et al. (1997; hereafter NFW) profile, over a very large radial range, and down to the center. Among other properties, for the JJ models the solution of the Jeans equations and the relevant global quantities entering the Virial Theorem can be expressed analytically. Therefore, the JJ models offer the {\it unique} opportunity to have a simple yet realistic family of galaxy models with a central MBH, allowing both for the fully analytical solution of the Bondi (isothermal) accretion problem {\it and} for the fully analytical solution of the Jeans equations; all this permits then a simple joint study of stellar dynamics and fluidodynamics without resorting to ad-hoc numerical codes. This paper is organized as follows. In Section 2 we recall the main properties of the Jaffe isothermal accretion solution, and in Section 3 we list the main properties of the JJ models. In Section 4 we show how the structural and dynamical properties of the stellar and dark matter components can be linked to the parameters appearing in the accretion solution. In Section 5 we examine the departure of the estimate of the mass accretion rate from the true value, when the estimate is obtained using as boundary values for the density and temperature those at points along the solution, at finite distance from the MBH. The main conclusions are summarized in Section 6. | The classical Bondi accretion theory is the tool commonly adopted in many investigations where an estimate of the accretion radius and the mass accretion rate is needed. In this paper, extending the results of previous works (KCP16, CP17), we focus on the case of isothermal accretion in two-component galaxies with a central MBH, and with radiation pressure contributed by electron scattering in the optically thin regime. In CP17 it was shown that the radial profile of the Mach number, and the critical eigenvalue of the isothermal accretion problem, can be expressed analitycally in Jaffe and Hernquist potentials with a central MBH. Here we adopt the two-component JJ galaxy models presented in CZ18. These are made of a Jaffe stellar component plus a dark halo such that the total density is also described by a Jaffe profile; all the relevant dynamical properties of JJ models, including the solution of the Jeans equations for the stellar component, can be expressed analytically. Therefore, the results of CP17 and CZ18 give the opportunity of building a family of two-component galaxy models where all the accretion properties can be given analytically, and then explored in detail, with no need to resort to numerical studies. The main results of this work can be summarized as follows. 1) The parameters describing accretion in the hydrodynamical solution of CP17 ($\MR$ and $\xi$) have been linked to the galaxy structure. In particular, it is assumed that the isothermal gas has a temperature $\Tinf$ proportional to the virial temperature of the stellar component, $\Tv$. Then, simple formulas are derived relating the galactic properties (as the effective radius, $\reff$, and the radius of influence of the MBH, $\Rinf$) with those describing accretion (as the Bondi radius $\rb$, and the sonic radius $\rmin$). The critical accretion parameter $\lambdat$ is also expressed as a function of the galactic properties. 2) For realistic galaxy structures, $\rb$ is of the order of a few$\times 10^{-3}\rs$, and $\Rinf$ is of the order of $\simeq 0.1\rb$. For $\Tinf =\Tv$, the sonic radius $\rmin$ is of the order of a few $\reff$. For moderately higher values of $\Tinf$, $\rmin$ suddenly drops to radii within $\rb$. The same happens also for a small increase of the polytropic index above unity, and this behavior is reminiscent of the similar jump shown by $\rmin$ in Hernquist models, as discussed in CP17. As a consequence, accretion in JJ models can switch from being supersonic over almost the whole galaxy to being everywhere subsonic, except for $r\la\rb$. An explanation for this phenomenon is given. 3) As for the isothermal accretion in one-component Jaffe models, the determination of the critical accretion parameter involves the solution of a quadratic equation, and there is only one sonic point for any choice of the parameters describing the galaxy. In presence of the galaxy, $\lambdat$ is several orders of magnitude larger than without the galaxy. It is found that Bondi accretion in JJ models in absence of a central MBH (or when $\chi =0$) is possible, provided that $\Tinf$ is lower than a critical value and we derive the esplicit formula for it. This critical value depends only on $\csig$, and is in the range $3/2\leq\Tinf/\Tv\leq 3$. It also determines the the jump in $\rmin$ in models with the central MBH. 4) We provide a few examples of accretion in realistic galaxy models, and present the resulting Mach number profiles, the trends of the accretion velocity and of the isotropic stellar velocity dispersion profiles. 5) We finally examine the problem of the deviation from the true value $\Mdott$ of an estimate of the mass accretion rate $\Mdotbe(r)$ obtained adopting the classical Bondi solution for accretion on a MBH, where the gas density and temperature at some finite distance from the center are inserted, as proxies for their values at infinity. The size of the departure of $\Mdotbe(r)$ from $\Mdott$, that is determined by the presence of the galaxy, is given as a function of the distance $r$ from the center. $\Mdotbe(r)$ {\it overestimates} $\Mdott$, if the gas density is taken in the very central regions, and {\it underestimates} $\Mdott$ if it is taken outside a few Bondi radii. This shows how sensitive to the model parameters is the determination of a physically based value for the so-called ``boost factor'' adopted in simulations, and that in general a universally valid prescription is impossible. | 18 | 8 | 1808.00871 |
1808 | 1808.08956_arXiv.txt | {We provide precise predictions for the hard photon spectrum resulting from neutral SU$(2)_W$ triplet (wino) dark matter annihilation. Our calculation is performed utilizing an effective field theory expansion around the endpoint region where the photon energy is near the wino mass. This has direct relevance to line searches at indirect detection experiments. We compute the spectrum at next-to-leading logarithmic (NLL) accuracy within the framework established by a factorization formula derived previously by our collaboration. This allows simultaneous resummation of large Sudakov logarithms (arising from a restricted final state) and Sommerfeld effects. Resummation at NLL accuracy shows good convergence of the perturbative series due to the smallness of the electroweak coupling constant -- scale variation yields uncertainties on our NLL prediction at the level of $5\%$. We highlight a number of interesting field theory effects that appear at NLL associated with the presence of electroweak symmetry breaking, which should have more general applicability. We also study the importance of using the full spectrum as compared with a single endpoint bin approximation when computing experimental limits. Our calculation provides a state of the art prediction for the hard photon spectrum that can be easily generalized to other DM candidates, allowing for the robust interpretation of data collected by current and future indirect detection experiments.} | Indirect detection is critical to the hunt for multi-TeV WIMP dark matter (DM). New data are continually being collected by current experiments, \emph{e.g.}~H.E.S.S.~\cite{Hinton:2004eu,Abramowski:2013ax,Abdallah:2018qtu}, HAWC~\cite{Sinnis:2004je,Harding:2015bua,Pretz:2015zja}, VERITAS~\cite{Weekes:2001pd,Holder:2006gi,Geringer-Sameth:2013cxy}, and MAGIC~\cite{FlixMolina:2005hv,Ahnen:2016qkx}, and a number of dedicated line searches for photons have been performed~\cite{Abramowski:2011hc,Abdallah:2018qtu}. Future experiments such as CTA~\cite{Consortium:2010bc,Acharya:2017ttl} will provide even greater sensitivity. Deriving the experimental ramifications these data will have on the parameter space of DM models requires precise predictions for the hard photon spectrum. Due to finite resolution effects inherent to the relevant experiments, a reliable prediction for not only the rate but also the shape of the spectrum is required to derive robust comparisons between theory and experiment~\cite{Baumgart:2017nsr}. The annihilation of TeV-scale DM is a multi-scale problem which is amenable to the application of effective field theory (EFT) techniques. In particular, non-relativistic EFTs can be used to treat the annihilating DM, and Soft-Collinear Effective Theory (SCET)~\cite{Bauer:2000ew,Bauer:2000yr, Bauer:2001ct, Bauer:2001yt} can be used for the final state radiation. The combination of these two EFTs~\cite{Baumgart:2014vma,Bauer:2014ula,Ovanesyan:2014fwa,Baumgart:2014saa,Baumgart:2015bpa,Ovanesyan:2016vkk} allows for the simultaneous resummation of Sudakov logarithms $\aW^n\log^{m}\big(M_\chi/\mW\big)$, with $m\leq 2\,n-1$ in the differential spectrum~\cite{Hryczuk:2011vi}, and Sommerfeld enhancement effects $\big( \aW M_\chi/\mW \big )^k$~\cite{Hisano:2003ec,Hisano:2004ds,Cirelli:2007xd,ArkaniHamed:2008qn,Blum:2016nrz}. In~\cite{Baumgart:2017nsr} we extended these EFT approaches to allow for the calculation of the hard photon spectrum in the endpoint region, where the photon energy $E_\gamma$ is near the DM mass $M_\chi$, as is relevant for line searches. Our framework additionally allows for the resummation of resolution effects $\aW^n\log^m\big(1-z\big)$ with $m\leq 2\,n-1$, where $z= E_{\gamma}/M_{\chi}$. Such logarithms are directly related to the experimental energy resolution, since $z$ quantifies the distance from the exclusive case, given by a line at $z=1$. A finite experimental resolution smears photons with a small $1-z$ into the expected exclusive event rate, and our calculation is able to realistically incorporate such effects. \begin{figure} \begin{center} \includegraphics[width=0.65\columnwidth]{figures/CumuSpec-Thermal} \end{center} \vspace{-15pt} \caption{The cross section for a thermal wino, \emph{i.e.}, with mass $M_\chi=2.9$ TeV~\cite{Beneke:2016ync}, as a function of resolution parameter $z_{\rm cut}$, showing our results at both LL and NLL. The NLL calculation significantly reduces uncertainties as compared to LL. A region appropriate for the H.E.S.S. experimental resolution, which is $\sim10\%$ at these energies~\cite{deNaurois:2009ud}, is shown in the grey band. This band is representative of the range of values that will contribute when our spectrum is convolved with the H.E.S.S. energy resolution function. } \label{fig:intro_fig} \end{figure} In this paper we extend this result to next-to-leading logarithmic (NLL) accuracy. To understand the importance of including the NLL corrections, we note that for the situation of interest here the logarithms $L$ become large enough so that $L^2\sim 1/\aW$. A leading logarithmic (LL) calculation then captures all terms scaling as $1$, and so should provide a good description of the shape of the distribution. However, a LL calculation does not probe higher order radiative corrections, and therefore typically has large uncertainties. On the other hand, an NLL calculation captures the first radiative corrections scaling like $\aW L$, and therefore typically provides a large reduction of the theoretical uncertainties. This reduction of uncertainties is clearly illustrated in \Fig{fig:intro_fig}, which shows a comparison of our earlier LL calculation with the NLL result achieved here. With NLL accuracy, the theory uncertainties become a subdominant contribution to the total uncertainty relevant experimentally. While our calculational framework is generally applicable to any heavy WIMP candidate, we will often specify to the case where the DM candidate is a wino, the neutral component of a triplet of SU$(2)_W$ with zero hypercharge. In addition to allowing us to illustrate our formalism, the wino is well motivated phenomenologically (see \emph{e.g.}~\cite{Giudice:1998xp,Randall:1998uk,ArkaniHamed:2004fb, ArkaniHamed:2004yi, Giudice:2004tc, Wells:2004di, Pierce:2004mk, Arvanitaki:2012ps, ArkaniHamed:2012gw, Hall:2012zp}), making our results of interest to current experiments. In a companion paper~\cite{Rinchiuso:2018ajn} we have used these results in a realistic H.E.S.S. forecast analysis to study the impact of having a complete description of the shape of the photon spectrum for wino searches. Here we provide the details of our NLL calculation, and perform a numerical study demonstrating that the theoretical uncertainty is significantly reduced when compared with the LL result, achieving an uncertainty from higher order corrections at the level of $5\%$. The cumulative cross section for $z\geq \zcut$ is shown in \Fig{fig:intro_fig} for a wino with a mass of $2.9$ TeV, which corresponds to the case where the thermal relic density matches the measured one~\cite{Beneke:2016ync}. Here $z_{\rm cut}$ restricts the cross section by allowing only photons with $z\ge \zcut$, see \Eq{eq:cumulative_def_intro}. We additionally show a band depicting the approximate values of $\zcut$ that correspond to the H.E.S.S. energy resolution. As can be seen, our calculation significantly reduces errors associated with the particle physics component of the annihilation cross section, which allows for the robust interpretation of experimental results in terms of DM model and astrophysical parameters. We also study the importance of using the full spectrum as compared with a single endpoint bin approximation when computing experimental limits by performing a mock H.E.S.S. analysis, and using the results of our forecasted H.E.S.S.~limits~\cite{Rinchiuso:2018ajn}. We find that the use of the full spectrum near the endpoint is crucial to preserve the desired accuracy, emphasizing the importance of our EFT formalism. Although the primary goal of this work is to provide a precision prediction for heavy WIMP annihilation in the endpoint region, a number of interesting features of SCET with broken gauge symmetry that have not previously appeared in the literature arise in our NLL calculation, and we devote several sections to their discussion. An outline of this paper is as follows. In \Sec{sec:review} we review the structure of the factorization formula for the endpoint region derived in \cite{Baumgart:2017nsr}, and prove that it remains valid at NLL accuracy. In \Sec{sec:NLL} we discuss the necessary formalism for achieving resummation at NLL accuracy, give explicit results for all one-loop anomalous dimensions, and solve the relevant renormalization group (RG) equations. In \Sec{sec:answer} we present analytic results for the cumulative and differential spectra at NLL accuracy, and comment on some interesting aspects of their structure. Numerical results and a study of the related theoretical uncertainties are given in \Sec{sec:results}. In \Sec{sec:compare} we compare the use of the full spectrum with a single endpoint bin in a mock H.E.S.S. analysis and with our forecasted limits, emphasizing the importance of having a complete description of the shape of the photon spectrum in the endpoint region. We conclude in \Sec{sec:conc}. In App.~\ref{app:explicitxsec} we collect all ingredients required for the cumulative and differential spectra, which are otherwise scattered throughout the paper. Finally in App.~\ref{app:intbrem} we demonstrate that our result matches existing fixed order calculations in the appropriate limit, and briefly comment on how the internal bremsstrahlung contribution, which is widely discussed in the literature, is reproduced in our result. | \label{sec:conc} In this paper, we have extended the calculation of the hard photon spectrum for wino annihilation in the endpoint region, as is relevant for indirect detection experiments, to NLL accuracy. This calculation was performed using an EFT framework developed in~\cite{Baumgart:2017nsr}, which facilitates the factorization of distinct physical effects. In particular, our result includes both the resummation of Sudakov logarithms and the Sommerfeld enhancement. The theoretical uncertainties of our calculation are of the order of $5\%$, which is a significant reduction as compared with our earlier LL prediction. In particular, the theory uncertainties are now sufficiently under control so as to make a subdominant contribution to the total uncertainty relevant for experimental exploration. In the course of our calculation we encountered a number of interesting effects associated with electroweak radiation and the presence of electroweak charged initial and final states. For example, we found that the non-electroweak-gauge-singlet nature of the incoming and outgoing states led to a non-trivial remaining Glauber phase in the final cross section. We think this would be interesting to explore further in a more general context than that considered here. The H.E.S.S.~telescope has collected a large dataset of photons from the Galactic Center region, with an energy resolution of $\sim 10\%$, permitting sensitive searches for spectral features. Using both a mock H.E.S.S.~analysis and a detailed forecasting framework, we studied the importance of using the full photon spectrum computed in this paper as compared with a single endpoint bin approximation when computing experimental limits. We find that the mapping to an effective bin width is a highly non-trivial function of the DM mass. This emphasizes the importance of having theoretical control over the shape of the distribution in the endpoint region, and not simply the photon count in an endpoint bin, for deriving accurate limits from experimental data. With an understanding of our factorization formula at NLL accuracy, it is now straightforward to calculate the spectrum at this accuracy for other heavy WIMP candidates, such as the pure Higgsino, the mixed bino-wino-Higgsino, or the minimal DM quintuplet~\cite{Cirelli:2005uq, Cirelli:2007xd, Cirelli:2008id, Cirelli:2009uv, Cirelli:2015bda,Mitridate:2017izz}. This will allow for the robust theoretical interpretation of indirect detection constraints for these compelling DM candidates from the wealth of data at current and future experiments. | 18 | 8 | 1808.08956 |
1808 | 1808.08398_arXiv.txt | Previous X-ray spectral analysis has revealed an increasing number of AGNs with high accretion rates where an outflow with a mildly relativistic velocity originates from the inner accretion disk. Here we report the detection of a new ultra-fast outflow (UFO) with a velocity of $v_{\rm out}=0.319^{+0.005}_{-0.008}c$ in addition to a relativistic disk reflection component in a poorly studied NLS1 WKK~4438, based on archival \nustar and \suzaku observations. The spectra of both \suzaku and \nustar observations show an Fe~\textsc{xxvi} absorption feature and the \suzaku data also show evidence for an Ar~\textsc{xviii} with the same blueshift. A super-solar argon abundance ($Z^{\prime}_{\rm Ar}>6Z_{\odot}$) and a slight iron over-abundance ($Z^{\prime}_{\rm Fe}=2.6^{+1.9}_{-2.0}Z_{\odot}$) are found in our spectral modelling. Based on Monte-Carlo simulations, the detection of the UFO is estimated to be around at 3$\sigma$ significance. The fast wind most likely arises from a radius of $\geq20r_g$ away from the central black hole. The disk is accreting at a high Eddington ratio ($L_{\rm bol}=0.4-0.7L_{\rm Edd}$). The mass outflow rate of the UFO is comparable with the disk mass inflow rate ($\dot M_{\rm out}>30\%\dot M_{\rm in}$), assuming a maximum covering factor. The kinetic power of the wind might not be high enough to have influence in AGN feedback ($\dot E_{\rm wind}/L_{\rm bol}\approx 3-5\%$) due to a relatively small column density ($12^{+9}_{-4}\times10^{22}$~cm$^{-2}$). However note that both the inferred velocity and the column density could be lower limits owing to the low viewing angle ($i=23^{+3}_{-2}$$^{\circ}$). | \label{introduction} Recently there has been increasing evidence showing the presence of absorption line features above 7~keV in the X-ray band of various sources, including Active Galactic Nuclei (AGNs) \citep[e.g.][]{chartas02, chartas03,cappi06,tombesi10,tombesi14,tombesi16}. These absorption line features are commonly interpreted as blueshifted Fe \textsc{xxv} or Fe \textsc{xxvi} K absorption in a highly ionized environment ($\log(\xi/$erg\,cm\,s$^{-1}$)>3) and can occasionally correspond to a very large line-of-sight outflow velocity of up to $0.2-0.4c$ \citep[e.g.][]{tombesi10}. Such outflows are often referred to as ultra-fast outflows (UFOs), and lie in the mildly relativistic regime indicating that the outflow is driven from the inner accretion disk. There are two key technical challenges when it comes to searching for ultra-fast outflows. The first is that they are close to the upper limit of the instrumental effective area of most soft energy cameras, such as \suzaku XIS and \xmm EPIC. Elements lighter than iron are generally fully ionized and therefore show weak or absent absorption features in the soft energy band. Only strong iron absorption feature remains above 7~keV where the signal-to-noise (S/N) and the spectral resolution are worse than in the soft band. The second is that the broadband continuum needs to be correctly modeled in order robustly determine the key absorption parameters \citep[e.g.][]{zoghbi15}. This is particularly critical for AGN sources, which often exhibit complex X-ray spectra with strong reflection. One popular theory for the origin of these extreme outflows is that the radiation pressure due to a high accretion rate drives the UFO (e.g. PDS~456; \citealt[][]{matzeu17}). It is therefore interesting to note the discovery of UFOs in ultraluminous X-ray sources \citep[ULXs; e.g.][]{pinto16, walton16, kosec18}, which appear to be sources accreting above their Eddington limit. Another ideal population to test this theory are Narrow Line Seyfert 1 (NLS1) galaxies. NLS1s are characterized by having low-mass, high-accretion-rate black holes in the center. For example, IRAS~13224$-$3809 accretes around the Eddington limit and shows a flux-dependent blueshifted Fe absorption feature above 8~keV \citep[e.g.][]{parker17a, parker17b,pinto17tmp,jiang18_tmp} which is interpreted as a UFO with velocity up to $0.236\pm0.006c$. WKK~4438 is a nearby (z=0.016) NLS1 galaxy hosting a low-mass supermassive black hole $M_{\rm BH}=2\times10^{6}M_{\odot}$ \citep[measured by the H$\beta$ line width in the optical band,][]{malizia08}. In this work, we analyse its archival X-ray spectra obtained by \suzaku and \nustar satellites, which show blueshifted Fe~\textsc{xxvi} and Ar~\textsc{xviii} absorption lines in addition to a relativistic reflection component. | \label{discussion} Archival \suzaku and \nustar data have revealed the NLS1 WKK~4438 shows evidence for both relativistic disk reflection and an ultra-fast outflow. The best-fit UFO parameters are a column density of $N^{\prime}_{\rm H}=12^{+9}_{-4}\times10^{22}$\,cm$^{-2}$, an ionization state of $\log(\xi^{\prime}$/erg\,cm\,s$^{-1})=3.9^{+0.4}_{-0.3}$, a slight iron over-abundance ($2.6^{+1.9}_{-2.0}Z_{\odot}$), an argon over-abundance (>6$Z_{\odot}$), and an outflowing velocity of $v_{\rm out}=0.319^{+0.005}_{-0.008}c$. The UFO absorption features are consistent when fitting the continuum only with the distant reflection model \texttt{xillver}. The inferred line significance is however higher with a distant reflection modelling. We discuss the physics properties of the outflow, other systematic uncertainties of the measurements, and future work in this section. To calculate the disk accretion rate, we apply an average bolometric correction factor $\kappa=20$ \citep{bolometric} to the 2--10\,keV flux in the \suzaku and \nustar observation. The bolometric luminosity of WKK~4438 is estimated to be $L_{\rm bol}=1.6-2.8\times10^{44}$\,erg\,s$^{-1}=0.4-0.7L_{\rm Edd}$, assuming $M_{\rm BH}=2\times10^6$\,$M_{\odot}$ \citep{malizia08}. A combination of high accretion rate and UFO is seen in other sources, such as rapidly accreting AGNs with high accretion rate (e.g. IRAS~13224$-$3809: \cite{parker17a,parker17b,pinto17tmp}; 1H0707$-$495: \cite{dauser12}; PDS~456 \cite{nardini15,reeves18_2}; PG~1211$+$143: \cite{fukumura15,reeves18}), and ULXs \citep{pinto16,walton16,kosec18}. The Eddington accretion rate for a black hole of $M_{\rm BH}=2\times10^6$\,$M_{\odot}$ is $9\times10^{20}$\,kg\,s$^{-1}$. Therefore the mass accretion rate of WKK~4438 is $\dot {M}_{\rm in}\approx0.006-0.01M_{\odot}$\,yr$^{-1}$. We estimate the mass outflow rate by following \citet{tombesi17}. A lower limit on the location of the wind can be derived from $r=2GM_{\rm BH}/v_{\rm out}^2\approx5.8\times10^{10}$\,m, which means that the wind is launched at $r\geq 20r_g$ away from the central SMBH. The mass outflow rate of the wind is $\dot {M}_{\rm out}=4 \pi \mu m_{p} r N_{\rm H} C_{\rm F} v_{\rm out}$, where $\mu=1.4$ is the average atomic mass per proton, $C_{\rm F}$ is the covering factor of the wind, and $m_{p}$ is the proton mass. Therefore, the lower limit on the mass outflow rate is $\dot {M}_{\rm out}\approx0.003C_{\rm F}M_{\odot}$yr$^{-1}$. The covering factor $C_{\rm F}$ of the wind remains unknown. The mass outflow rate is comparable with the mass accretion rate, assuming a maximum covering factor ($C_{\rm F}\approx1$ for PDS~456 \citep{nardini15} and IRAS~F11119$+$3257 \citep{tombesi15}). The kinetic power of the wind is then $\dot E_{\rm wind}=(1/2)\dot {M}_{\rm out}v_{\rm out}^2\approx9.59\times10^{42}$~erg\,s$^{-1}\approx3-5\%L_{\rm bol}$, assuming maximum covering factor. The low $\dot E_{\rm wind}/L_{\rm bol}$ ratio is due to a small column density of the wind and indicates the energetics might not be high enough to have an influence in AGN feedback \citep[e.g.][]{dimatteo05}. However, note that both the inferred velocity and the column density could be lower limits owing to the low inclination we infer, assuming the wind still has a roughly equatorial geometry. In order to understand the high argon abundance required by the best-fit spectral model, we first investigated relative argon abundance against several other elements, such as silicon, sulfur, calcium, iron and oxygen, by fitting the spectra with the \texttt{warmabs} model. It turns out that the spectral fitting requires a high $Z^{\prime}_{\rm{Ar}}/Z^{\prime}_{\rm{Fe}}>4.2$ in the outflow. When the UFO and disk iron abundances are linked during the fit ($Z^\prime_{\rm Fe}=Z_{\rm Fe}$), a solar iron abundance is required ($Z^{\prime}_{\rm Fe}=Z_{\rm Fe}=0.96^{+0.66}_{-0.18}Z_{\odot}$) with $\chi^2=1364.1$, slightly worse than the fit with $Z^{\prime}_{\rm Fe}$ and $Z_{\rm Fe}$ both as a free parameter. The other parameter values do not change too much when the two iron abundances are linked. Second, we also tried fitting the data with a single matallicity ($Z^{\prime}$), with the abundance of all elements heavier than He linked together. It offers a metallicity value of $Z^{\prime}>0.6Z_{\odot}$ with $\chi^2=1388.4$. Finally, if we allow for a free iron abundance in addition to a single metallicity for the rest of the heavy elements, we obtained a fit with $\chi^2=1381.2$ ($Z^\prime_{\rm Fe}<3Z_{\odot}$ and $Z^{\prime}>0.6Z_{\odot}$) and only the 10~keV absorption feature being fitted. We note that \citet{tombesi10} report a similar scenario, where absorption from iron and argon with a common blueshift, and no other lines detected, for the Seyfert galaxy NGC~7582. \citet{dauser12} analysed the 1H0707-495 \xmm spectra and found a possible P Cygni profile of H-like Ar in the 2008 observations, where the redshift of the Ar line feature is however different from other absorption lines. In contrast, another NLS1 IRAS~13224$-$3809 shows a series of absorption lines in the middle energy band, including Ne~\textsc{x}, S~\textsc{xvi} and Si~\textsc{xiv} absorption lines, but shows no evidence of Ar~\textsc{xviii} absorption line in the spectra \citep{jiang18_tmp}. Although, the variability spectrum of IRAS~13224$-$3809 shows some evidence of Ar~\textsc{xviii} feature \citep{parker17b}. By fitting the spectra with the relativistic reflection model \texttt{relxilllp}, we obtain a disk viewing angle of $i=23^{+3}_{-2}$$^{\circ}$. Such a small viewing angle is unusual for a source with visible UFO absorption features due to the opening angle of the wind. Other sources where UFO is detected show evidence of large disk viewing angle when being fitted with relativistic disk reflection model, such as 1H0707$-$495 \citep[$i\approx50$$^{\circ}$,][]{dauser12}, PDS~456 \citep[$i=65\pm2$$^{\circ}$,][]{chiang17}, IRAS~13224$-$3809 \citep[$i=67\pm3$$^{\circ}$,][]{jiang18_tmp}. An alternative explanation of the absorption features is ionized materials corotating above the disk, where the relativistic velocities occur naturally \citep{gallo11}. Such a model has been successfully applied to PG~1211$+$143 \citep{gallo13} and IRAS~13224$-$3809 (Fabian et al. in prep). However, this scenario seems to be unlikely for WKK~4438 due to its small viewing angle and truncated disk. It is now becoming apparent that these extreme outflows are variable phenomena in various sources, such as IRAS~13224$-$3809 \citep{parker17a,pinto17tmp}, Mrk~509 \citep{cappi09}, PDS~456 \citep{reeves09,nardini15,matzeu1}, and PG1211$+$143 \citep{pounds03,reeves18}. In many of these cases, the black hole masses are high (e.g. $M_{\rm BH}\approx10^9$M$_{\odot}$ for PDS~456 and $M_{\rm BH}\approx10^8$M$_{\odot}$ for PG1211$+$143). However, similar to IRAS~13224$-$3809, the mass of the black hole in WKK~4438 is rather low ($M_{\rm BH}\approx2\times10^6$\,$M_{\odot}$, \citet{malizia08}). This means the outflow in WKK~4438 is potentially of particular interest, as variability timescales are generally expected to scale with black hole mass. It may therefore be possible to study the variability of the outflow in WKK~4438 -- and any potential response to intrinsic changes in the source -- with a few deep observations in the future. Unfortunately, the data currently available do not have sufficient signal-to-noise or broadband coverage to undertake more studies at the present time. | 18 | 8 | 1808.08398 |
1808 | 1808.01793_arXiv.txt | {} {With our low-resolution spectroscopic observing program for selected blue proper motion stars, we tried to find new white dwarfs (WDs) in the solar neighbourhood.} {We used the L{\'e}pine \& Shara Proper Motion (LSPM) catalogue with a lower proper motion limit of 150\,mas/yr and the Second US Naval Observatory CCD Astrograph Catalog (UCAC2) for proper motions down to about 90\,mas/yr. The LSPM and UCAC2 photometry was combined with Two Micron All Sky Survey (2MASS) near-infrared (NIR) photometry. Targets selected according to their blue optical-to-NIR and NIR colours were observed mainly at Calar Alto. The spectra were classified by comparison with a large number of already known comparison objects, including WDs, simultaneously observed within our program. Gaia DR2 parallaxes and colours were used to confirm or reject spectroscopic WD candidates and to derive improved effective temperatures. } {We found ten new WDs at distances between 24.4\,pc and 79.8\,pc, including six hot DA WDs: \object{GD 221} (DA2.0), \object{HD 166435 B} (DA2.2), \object{GD 277} (DA2.2), \object{2MASS J19293865+1117523} (DA2.4), \object{2MASS J05280449+4105253} (DA3.6), and \object{2MASS J05005185-0930549} (DA4.2). The latter is rather bright ($G\approx12.6$) and with its Gaia DR2 parallax of $\approx14$\,mas it appears overluminous by about 3\,mag compared to the WD sequence in the Gaia DR2 colour-magnitude diagram. It may be the closest extremely low mass (ELM) WD to the Sun. We further classified \object{2MASS J07035743+2534184} as DB4.1. With its distance of 25.6\,pc it is the second nearest known representative of its class. With \object{GD 28} (DA6.1), \object{LP 740-47} (DA7.5), and \object{LSPM J1919+4527} (DC10.3) three additional cool WDs were found. Gaia DR2 parallaxes showed us that four of our candidates but also two previously supposed WDs (\object{WD 1004+665} and \object{LSPM J1445+2527}) are in fact distant Galactic halo stars with high tangential velocities. Among our rejected WD candidates, we identified a bright ($G=13.4$\,mag) G-type carbon dwarf, \object{LSPM J0937+2803}, at a distance of 272\,pc. } {} | \label{Sect_intro} The stellar content of the Solar neighbourhood provides important information for the study of the stellar initial mass function, Galactic structure and evolution. White dwarfs (WDs) play an important role in our understanding of these fundamental astrophysical issues \citepads{2005ARA&A..43..247R,2014ApJ...791...92T}. In particular, the cool WDs provide age estimates of the Galactic disk \citepads{2015MNRAS.449.3966G,2017ApJ...837..162K} and information on the baryonic dark matter in the Galaxy \citepads{2001ApJ...559..942R,2002A&A...395..779M,2003MNRAS.339..817F}. The nearest WDs are also excellent targets for divers methods to search for extrasolar planets \citepads{2010ApJ...708..411K,2015A&A...579L...8X}. As the general census of nearby stars has been slowly improved in the pre-Gaia era \citepads[see e.g.][]{2018AJ....155..265H} the sample of known WDs in the solar neighbourhood experienced a similar evolution. While the WD census was believed to be complete within 13\,pc \citepads{2016MNRAS.462.2295H}, during the last ten years many efforts were undertaken to identify more WDs in the extended solar neighbourhood (e.g. for the 25\,pc and 40\,pc samples) by measuring their parallaxes \citepads{2009AJ....137.4547S,2017AJ....154...32S} and using spectroscopic surveys \citepads{2011ApJ...743..138G,2012MNRAS.425.1394K, 2013AJ....145..136L,2014AJ....147..129S}. The physical properties of nearby WDs were investigated in more detail by \citetads{2012ApJS..199...29G}, \citetads{2015ApJS..219...19L}, \citetads{2017MNRAS.465.2849T}, \citetads{2017MNRAS.467.4970H}, and \citetads{2017AJ....154..118S}. The overall number of spectroscopically classified WDs (including more distant objects) was increased fivefold with the catalogues of \citetads{2013ApJS..204....5K} and \citetads{2015MNRAS.446.4078K,2016MNRAS.455.3413K}, based on spectroscopic data from the Sloan Digital Sky Survey \citepads[SDSS;][]{2000AJ....120.1579Y}. Despite of all the above mentioned efforts, even the generally assumed completeness of the 13\,pc WD sample was shown to be not correct, as the recent discovery of the cool WD companion of a new nearby red dwarf star \object{TYC 3980-1081-1 B} \citepads{2018A&A...613A..26S} at a distance of 8.3\,pc indicated. This discovery was based on astrometric data from Gaia DR1 \citepads{2016A&A...595A...4L}, the 5th United States Naval Observatory CCD Astrograph Catalog \citepads[UCAC5;][]{2017AJ....153..166Z}, and the U.S. Naval Observatory Robotic Astrometric Telescope \citepads[URAT;][]{2015AJ....150..101Z} Parallax Catalog \citepads[UPC;][]{2016yCat.1333....0F}. The WD status of \object{TYC 3980-1081-1 B} is now confirmed by its Gaia DR2 \citepads{2018arXiv180409365G} parallax of 118.12$\pm$0.02\,mas. Most searches for previously unrecognised nearby stars in the pre-Gaia era investigated targets selected from combined colour- and proper motion surveys. In 2008, we started a spectroscopic follow-up programme to search for missing WDs in the solar neighbourhood in a sample of blue proper motion stars. The L{\'e}pine \& Shara Proper Motion \citepads[LSPM; ][]{2005AJ....129.1483L} catalogue of the northern sky and the Second US Naval Observatory CCD Astrograph Catalog \citepads[UCAC2; ][]{2004AJ....127.3043Z} covering the sky area from $-90\degr$ to $+40\degr$ declination (going up to $+52\degr$ in some areas), and the Two Micron All Sky Survey \citepads[2MASS;][]{2006AJ....131.1163S} providing accurate near-infrared (NIR) photometry served as our main input data with respect to proper motions and photometry. The results of our WD survey were not published so far, except for one new WD (\object{HD 166435 B} = \object{CA376}) that was independently discovered in a search for common proper motion WD companions of known nearby stars \citepads{2018A&A...613A..26S}. This paper is organised as follows: In Sect.~\ref{Sect_sample} we describe the selection of our target stars for the spectroscopic observations, which were mainly carried out at the Calar Alto observatory (Sect.~\ref{Sect_spec}). Sect.~\ref{Sect_class} deals with our spectroscopic classification of WD candidates of different classes, whereas Sects.~\ref{Sect_Plxg} and \ref{Sect_BP_RP} show how well our results are confirmed by Gaia DR2 parallaxes and colours. Finally, we review previous investigations on our confirmed and rejected WDs and briefly discuss the properties of the most interesting objects in Sect.~\ref{Sect_discuss}. | \label{Sect_discuss} \subsection{Classification and kinematics of rejected WDs} \label{Sect_notes_rejected} Three objects (\object{CA171}, \object{CA308}, and \object{CA351}) out of six rejected WD candidates and known WDs described in Sects.~\ref{Sect_DQAZ} and \ref{Sect_rejected} were photometrically classified as FGK stars by \citetads{2010PASP..122.1437P}, and one object (\object{CA192}) was listed as K star candidate in the catalogue of \citetads{2001KFNT...17..409K}. The WD classification of \object{CA297} was rejected on the basis of an SDSS M0V spectrum already noted in Sect.~\ref{Sect_rejected}. The largest tangential velocity among our new WDs is 77\,km/s for \object{CA399}, possibly indicating its membership in the Galactic thick disk population. On the other hand, all rejected WDs exhibit tangential velocities between $\approx210$\,km/s and $\approx520$\,km/s, relating them to the Galactic halo. They are all located at high Galactic latitudes (between $+40\degr$ and $+71\degr$), whereas the ten new WDs lie near the Galactic plane between $-29\degr$ and $+40\degr$. The Gaia parallax of \object{CA171} (Table~\ref{Tab_6nonWDs}) confirms its classification as a G-type carbon dwarf. With the corresponding distance of $\approx272$\,pc it is a relatively nearby representative of this class of objects. The vast majority of G-type dC stars presented by \citetads{2013ApJ...765...12G} are much fainter than \object{CA171} and show typical total proper motions below 30\,mas/yr (see his fig.~11). Compared to other nearby carbon dwarfs \citepads{2018AJ....155..252H}, distance and tangential velocity of \object{CA171} are typical for the Galactic halo population. But, compared to other carbon dwarfs it appears much bluer ($J-K_{\rm s}\approx+0.4$\,mag) and relatively bright ($J\approx12.1$\,mag) not only in the NIR (2MASS) but also in the optical ($B=14.59$\,mag, $V=13.55$\,mag) \citepads[APASS;][]{2016yCat.2336....0H}. \subsection{Previous classification of new WDs} \label{Sect_prev_class} Three of the new WDs, \object{CA037} (= \object{GD 277}), \object{CA054} (= \object{GD 28}), and \object{CA399} (= \object{GD 221}), were already considered as WD candidates by \citetads{1980LowOB...8..157G}. Later on two of them were classified as non-WDs in \citetads{2010PASP..122.1437P} based on multi-colour photometry (\object{CA037} as A7III candidate and \object{CA054} as G8V candidate). Another new WD, \object{CA397}, was previously listed as a \textit{Kepler} target star with estimated 1.5 solar masses and about solar metallicity \citepads{2014ApJS..211....2H}. Two of the new WDs, \object{CA054} and \object{CA305} were listed in Luyten's WD catalogue \citepads{1999yCat.3070....0L}, both with a spectral or colour class ``f''. \subsection{URAT parallaxes of new WDs} \label{Sect_UPC} For four of the new WDs (\object{CA037}, \object{CA103}, \object{CA376}, and \object{CA397}) first trigonometric parallaxes of lower precision were already published before Gaia DR2 in the URAT \citepads{2015AJ....150..101Z} Parallax Catalog \citepads[UPC;][]{2016yCat.1333....0F,2018yCat.1344....0F} described by \citetads{2016arXiv160406739F,2018AJ....155..176F}. For two of them, \object{CA037} and \object{CA103}, the UPC parallaxes are in excellent agreement with the Gaia DR2 values, and for \object{CA397} the UPC parallax agrees within the typical UPC error bars ($\pm$3-6\,mas) with the Gaia DR2 parallax. Only for \object{CA376}, the UPC parallax is about 10\,mas larger than the Gaia DR2 parallax. \subsection{Wide binary companions of our new WDs} \label{Sect_wide_binaries} Wide binaries were traditionally found among common proper motion stars in high proper motion catalogues. Now, in the Gaia era, we can use very accurate measurements of both proper motions and parallaxes for an almost complete sample of stars in the solar neighbourhood. Therefore, we checked our ten new WDs for stars with common proper motion and parallax in Gaia DR2 data. We used a search radius of 3600\,arcsec and selected only Gaia DR2 sources with very similar parallaxes and then compared their proper motions. In addition to the already known G-type primary of \object{CA376} \citepads[= \object{HD 166435 B}; see][]{2018A&A...613A..26S} with an angular separation of about 29.2\,arcsec, we found \object{2MASS J19293859+1118050} a red ($BP-RP\approx+2.62$\,mag) companion of \object{CA400} separated by only about 12.7\,arcsec. Its parallax (27.38$\pm$0.07\,mas) is in perfect agreement with that of \object{CA400}, whereas its proper motion of (+18.76$\pm$0.11, 99.80$\pm$0.09)\,mas/yr is only slightly deviating (Table~\ref{Tab_10WDs}), as expected due to orbital motion \citepads{2018A&A...613A..26S}. We derived an absolute $J$ magnitude of $\approx$7.3\,mag for this red companion from its 2MASS photometry and Gaia DR2 parallax corresponding to a spectral type of $\approx$M3 according to the relationship between absolute magnitudes and spectral types \citepads{2005A&A...442..211S}. \subsection{One of the nearest DB WDs: CA103} \label{Sect_CA103} With a distance of only $\approx$26.60\,pc, \object{CA103} is the second nearest of our new WDs after \object{CA376} at $\approx$24.45\,pc (based on Gaia DR2 parallaxes). There were only two DB WDs in the 25\,pc WD sample \citepads{2016MNRAS.462.2295H}, and one of these two, \object{WD 2058+342}, is now found to be at a distance of $\approx$52.51\,pc according to its Gaia DR2 parallax.\footnote{Investigating the Gaia DR2 20\,pc sample, \citetads{2018arXiv180512590H} idendified 128 known and 11 new WDs, whereas 57 (!) former members of the 20\,pc sample were found to be located at larger distances.} Therefore, \object{CA103} may in fact be the second nearest of all DB WDs after \object{WD 1917-074} at $\approx$10.50\,pc. Among our new WDs (Table~\ref{Tab_10WDs}), \object{CA103} has a very small proper motion for its distance, and consequently the smallest tangential velocity of $\approx$10\,km/s, indicating a young age. Interestingly, \object{CA103} is listed in the K2 Ecliptic Plane Input Catalog \citepads[EPIC;][]{2017yCat.4034....0H}. It represents a relatively bright ($G\approx13.9$\,mag, $J\approx14.0$\,mag) target for various follow-up observations. \subsection{CA078, a nearby extremely low mass WD candidate} \label{Sect_CA078} The multiplicity of nearby WDs, in particular double WDs (DWDs) and WD main-sequence binaries (WDMS), was investigated by \citetads{2017A&A...602A..16T}. They listed two unresolved DWDs and eight unresolved DWD candidates within 25\,pc (mostly DA types between DA5.6 and DA9.9) in their Table 1. In addition, one of their only four resolved DWDs (DC10) has a small angular separation of 1.4\,arcsec. They also compared the observed multiplicity of WDs with star formation and evolution model predictions and found a discrepancy for resolved DWDs, more than ten of these are apparently missing within 20\,pc. \begin{figure} \centering \includegraphics[width=7.8cm]{aa33700_f12.jpg} \caption{Balmer jump as a function of effective temperature measured in low-gravity model spectra (small filled squares connected by solid lines) between $\log{g}=5.0$ (top) and $\log{g}=6.5$ (bottom). The dashed vertical line indicates the effective temperature of \object{CA078} (Sect.~\ref{Sect_BP_RP}). Dashed and dotted horizontal lines indicate our observed value for \object{CA078} and our estimated measurement uncertainty. } \label{Fig_Bjump} \end{figure} Our overluminous WD \object{CA078} lies at a much larger distance of about 71.6\,pc, where the chances to resolve a possible multiple system are lower, and has a slightly earlier spectral type (DA4.2$\pm$0.5) than the above mentioned unresolved DWDs. In addition, an equal-mass WD binary would appear only 0.75\,mag brighter than a single WD. As already mentioned in Sect.~\ref{Sect_Plxg}, \object{CA078} belongs to the high-quality Gaia DR2 100\,pc sample and its location about 3\,mag above the WD sequence in the optical-to-NIR CMD (Fig.~\ref{Fig_MG_G_J}) and optical CMD (Fig.~\ref{Fig_MG_BP_RP}) requires an alternative explanation different from simple binarity. An elusive class of WDs with low surface gravities of $5\lesssim \log{g} \lesssim7$ and effective temperatures in the range of 8000\,K$\lesssim T_{\rm eff} \lesssim$22\,000\,K are the so-called extremely low mass (ELM) WDs \citepads{2016ApJ...818..155B}. With our estimated effective temperature of about 11\,800\,K well within this interval, can the overluminosity of \object{CA078} be explained with its possible ELM status? Our Balmer line equivalent width measurements (Figs.~\ref{Fig_HaHbEW} and \ref{Fig_HbHgEW}), in particular of the well-measured H$\beta$ and H$\gamma$ lines, already provided a hint on the possible low gravity ($\log{g}\lesssim7$) of \object{CA078} (Sect.~\ref{Sect_modelspec}). All known DA3.6-DA4.5 shown (in Figs.~\ref{Fig_HaHbEW} and \ref{Fig_HbHgEW} have much larger H$\beta$ (and H$\alpha$) line widths than \object{CA078}, whereas the H$\gamma$ line widths are similar. In addition to the Balmer line widths, we can also measure the gravity-dependent Balmer jump. We defined it as the flux ratio between the mean spectral energy densities in the wavelength intervals 4200-4240\,{\AA} and 3700-3740\,{\AA}, which were covered both by the model and the observed spectra. Those flux ratios were also measured in spectra normalised to a continuum (instead of measuring in the original spectra) to avoid a colour term due to calibration uncertainties of the instrumental response function for our observed spectra. In Fig.~\ref{Fig_Bjump} we compare the size of the Balmer jump in the observed and the model spectra. The graph indicates a low value of about $\log{g}=6$ of the surface gravity of \object{CA078} with an upper limit of about $\log{g}=6.5$. This is consistent with an ELM status of \object{CA078}. \begin{figure} \centering \includegraphics[width=\hsize]{aa33700_f13.jpg} \caption{Previous distances of the known nearest ELM WDs (see text) and their new distances derived from Gaia DR2 parallaxes. The red dashed line represents equality. Objects shifted by Gaia to much larger distances are labelled and shown in black. Other objects are shown in red, the three nearest of which are labelled.} \label{Fig_ELM} \end{figure} Gaia DR2 and SDSS data were already combined by \citetads{2018arXiv180504070P} in a study of so-called sdA stars, including (pre-)ELM WDs of low gravity and large radius. Using the same method, we estimate a radius of $0.0429^{+0.0213}_{-0.0122} r_{\sun}$ and a large uncertain distance of $913^{+282}_{-403}$\,pc for the known ELM \object{SDSS J091709.55+463821.7} \citepads{2016ApJ...818..155B}, which shows similar atmospheric parameters to \object{CA078} ($\log{g}\approx6$ and $T_{\rm eff}\approx$12\,000\,K). The estimated radius is about four times larger than the canonical WD radius corresponding to an increase in magnitude of about 3\,mag, similar to what we see for \object{CA078} in Figs.~\ref{Fig_MG_G_J} and \ref{Fig_MG_BP_RP}. While \object{SDSS J091709.55+463821.7} has only an uncertain Gaia DR2 parallax and does not fall in the high-quality Gaia DR2 100\,pc sample, we can compute its absolute $G$ magnitude as about 9.1\,mag, which is only slighly larger than that of \object{CA078} ($\approx$8.3\,mag). Both objects have almost exactly the same zero colour indices $BP-RP$ and would appear very close to each other in the Gaia DR2 CMD of Fig.~\ref{Fig_MG_BP_RP}. Therefore, we conclude that \object{CA078} is probably an ELM WD similar to \object{SDSS J091709.55+463821.7} but at much smaller distance. We have cross-identified the entire ELM survey sample of \citetads{2016ApJ...818..155B} with Gaia DR2. Out of all 88 objects, 18 (with previous distance estimates between 0.42\,kpc and 7.77\,kpc) have negative or no parallaxes measured in Gaia DR2. Out of the remaining 70 ELM WDs, there are only 14 previously considered as relatively nearby, with distances between 0.1\,kpc and 0.3\,kpc according to \citetads{2016ApJ...818..155B}. However, as can be seen in Fig.~\ref{Fig_ELM}, six of them (shown in black) are shifted to much larger distances according to their Gaia DR2 parallaxes, including the previously assumed nearest ELM WD \object{UCAC2 48250979}. On the other hand, eight objects (shown in red) have only slightly larger Gaia DR2 distances compared to their previous distance estimates. The three nearest new ELM WDs are \object{LP 413-40}, \object{SDSS J075519.48+480034.0}, and \object{HS 1102+0934} at distances of 182\,pc, 183\,pc, and 191\,pc. From these three objects, only \object{LP 413-40} is included in the clean sample of ELM WD binaries \citepads{2016ApJ...818..155B} and lies with $BP-RP\approx+0.45$\,mag and an absolute G magnitude of $\approx$10.34\,mag almost 3\,mag above the normal WD sequence in the Gaia CMD (cf. Fig.~\ref{Fig_MG_BP_RP}), similar to \object{CA078}. We conclude that \object{CA078}, which is located about three times closer to the Sun and has a higher effective temperature than \object{LP 413-40}, represents a good new candidate for the nearest ELM WD. \subsection{Outlook} \label{Sect_outlook} We have shown the Gaia DR2 colours $BP-RP$ to be a very effective tool for estimating effective temperatures of WDs. This can probably be further enhanced when a larger sample of known WDs will be considered. The polynomial fitting with our small sample of known WDs already allowed us to reach uncertainties between $\pm$1100\,K and $\pm$300\,K for hot and cool WDs. So, we can e.g. for the recently discovered new WD member of the 10\,pc sample, \object{TYC 3980-1081-1 B} \citepads{2018A&A...613A..26S}, estimate an effective temperature of 5100$\pm$300\,K based on its Gaia DR2 colour of $BP-RP\approx1.06$\,mag. We think that our new WDs, but also some of our rejected WDs, deserve further attention with follow-up observations, including higher-resolution spectroscopy, radial velocity monitoring, variability analysis, and imaging (to search for possible companions). In particular, the nature of the overluminous DA4.2$\pm$0.5 ELM WD candidate \object{CA078} needs to be clarified. If confirmed, this would be the nearest ELM WD. The DB4.1$\pm$0.4 WD \object{CA103} can be studied in more detail as one of the nearest representatives of its class. Among our rejected WDs, the G-type carbon dwarf \object{CA171} is probably a promising target for high-resolution spectroscopy, as it is as bright ($G\approx13.35$\,mag) as the well-known and metal-poor ([Fe/H]$=-4$) dC star \object{G 77-61}. The latter is according to \citetads{2018RNAAS...2b..43M} the only dC star with a detailed abundance analysis \citepads{2005A&A...434.1117P} so far. | 18 | 8 | 1808.01793 |
1808 | 1808.08035_arXiv.txt | Two EAS arrays during a day have recorded 3 particles with energies above 30 EeV arriving from the same sky region. Two events were registered by the Yakutsk array and one - by the Telescope Array. Two Yakutsk events were estimated to be same energy and the Telescope Array shower's energy was almost 2 times higher. This indicates the same magnetic rigidity of all three particles, if the charge differs by a factor of 2. The relatively short sequence of all three events and their "monochromaticity" in rigidity can be due to the magnetic separation of particles in the acceleration process and propagation of the beam. | 18 | 8 | 1808.08035 |
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1808 | 1808.06819_arXiv.txt | Proxima b is a terrestrial exoplanet orbiting in the habitable zone of our closest star Proxima Centauri. The separation between the planet and the star is about 40 mas and this is with current instruments only reachable with direct imaging, using a visual extreme AO system like SPHERE/ZIMPOL. Unfortunately, the planet falls under the first airy ring at 2$\lambda$/D in the I band, which degrades achievable contrast. We present the design, optical simulations and testing of an amplitude pupil mask for ZIMPOL that reshapes the PSF, increasing the contrast at $r = 2\lambda$/D about an order of magnitude. The simple mask can be inserted directly into the current setup of SPHERE. | \label{sec:intro} % The last decade has sparked a revolution in the field of exoplanet science with technologies and observing and data processing techniques that have led to the discovery of over three thousand planets. It is now suspected that at least every second star has one or more planets, and the diversity of the planet parameters of radii, mass and system architecture is astonishing. Several exo-earth candidates have been discovered, terrestrial planets which are located in the habitable zone of the star, have atmospheres and could have liquid water on their surface. None of these objects has been directly observed but rather found indirectly through transit and radial velocity methods. Direct imaging of exoplanets has been restricted to hot gas giants with masses a few times the mass of Jupiter. This is mainly due to the extremely challenging task of detecting faint signals, many orders of magnitude smaller than the star, which favors the discovery of self luminous young giant planets far away from the star. A recently discovered terrestrial planet of great interest is Proxima Centauri b (HIP 70890 b) \cite{Anglada-Escude2016}, which was discovered through radial velocity method. It is orbiting around the nearest star at a distance of 1.295 parsecs with an orbit of 11.2 days and a semi-major axis of 0.048 AU. At quadrature phase the angular distance is 37 milli-arcseconds (mas) and when observed with a 8 m-class telescope it corresponds to two diffracted beam widths (2 $\lambda$/D) in the I-band. Since the host star is a red dwarf Proxima b might have conditions that make it habitable \cite{Dong2017} \cite{Meadows2016}. In Figure \ref{exoplanets} of Lovis et al. (2016)\cite{Lovis2016} known extrasolar planets with small separations are plotted with the y-axis indicating the expected planet-to-star flux ratio. There are various limitations that prevent current instruments to detect these objects. First, diffraction limits the wavelength range due to the current aperture size of the largest telescopes and future Extremely Large Telescopes (ELTs) that would alleviate this limit are projected in 7 to 10 years. Secondly, the requirement in contrast is orders of magnitude smaller than what is currently achieved even with the "planet-hunter" instruments that operate mainly in the Near-Infrared (H-band) and are equipped with extreme adaptive optics and coronagraphs. Lovis et al. (2016) \cite{Lovis2016} have proposed a coupling of the ESSPRESSO spectrograph at VLT with the high contrast imaging path of SPHERE as a High Contrast High Resolution Spectroscopy configuration and observe Proxima b with the hope to extract a spectrum of its atmosphere. As the authors suggest in this paper it would be interesting to first observe it with ZIMPOL (Zurich IMaging POLarimeter) in polarimetry on SPHERE. \begin{figure} \begin{center} \includegraphics[scale=0.5]{Estimated-planet-to-star-contrast-in-reflected-light-for-known-exoplanets-as-a-function.png} \end{center} \caption{Expected planet to star flux ratio in reflected light versus angular separation. The color indicates the estimated temperature of the objects and the vertical lines correspond to integer multiples of the resolving power of an 8-m class telescope at 750 nm. A few of the interesting targets are indicated by their name, including Proxima b. Plot used with permission from Lovis et al. (2016) \cite{Lovis2016}.} \label{exoplanets} \end{figure} \subsection*{SPHERE/ZIMPOL instrument at VLT} ZIMPOL is a high contrast polarimeter designed to search for exoplanets in reflected light and is part of the SPHERE instrument\cite{schmid2018} at VLT. Operating in the visible to near-infrared wavelengths, with an extreme adaptive optics system \cite{fusco2014final} and using polarimetric imaging one can reach deep contrasts of up to $10^{-6}$ at close separations, whilst providing high angular resolution. Light of the star reflected off the planet's atmosphere results in a fraction of the light being polarized. ZIMPOL observes two orthogonal polarization directions simultaneously, and while the signal of the star can be found in both directions the polarized planet signal only ends up in one. One can therefore subtract the star signal efficiently with polarimetric differential imaging (PDI) and faint sources of scattered or reflected light can be detected. This technique has been highly successful in detecting and characterizing faint circumstellar disks\cite{engler2017hip}. An ongoing search of planets in reflected light with ZIMPOL provides for nearby bright stars planet contrast limits if $10^{-6}$ to $10^{-8}$ outside of 0.2" separation (Hunziker et al, in preparation). An important distinction of the planets searched in reflected light is that objects close to their star are favored since they receive much more irradiation, which in turn leads to stronger emission of polarized light. For observing Proxima b the requirements are still at the upper limit of what is theoretically achievable with the current installment of ZIMPOL. The two biggest limitations for detecting signals at separations below 3 $\lambda$/D are some uncorrected instrumental polarization induced due to the beam shift effect as described in Schmid et al (2018)\cite{schmid2018}, and the residual light from the first airy ring. The first airy ring for an obstructed aperture such as the VLT still encompasses a non-negligible fraction of the host star PSF which prevents to reach deep contrasts. At the moment the achievable polarimetric contrast is 3 $10^{-7}$ at a separation of 0.2" and the goal of this study is to provide a simple solution that can be inserted into the setup of SPHERE and would provide a similar contrast at closer separation. The idea is to design an binary amplitude pupil mask that reshapes the pupil and improves contrast in the region of interest. | The prospect of directly imaging Proxima b is exciting and we believe it is something worth pursuing even with the current technology. We show that inserting an amplitude mask in the ZIMPOL instrument would improve the contrast and provide about an order of magnitude gain in contrast. This would be a low risk high gain approach without the need of an extensive upgrade of the instrument, implementing a simple and robust concept. Even a non detection of Proxima b would be interesting, giving some limits for atmospheric models and designs of future instrument upgrades. Furthermore a detection of an exoplanet in reflected light would be of great achievement for the ZIMPOL/SPHERE instrument and pave the way for similar instruments at the future ELTs. | 18 | 8 | 1808.06819 |
1808 | 1808.10472_arXiv.txt | V350 Sgr is a classical Cepheid suitable for mass determination. It has a hot companion which is prominent in the ultraviolet and which is not itself a binary. We have obtained two high resolution echelle spectra of the companion at orbital velocity maximum and minimum with the Hubble Space Telescope (HST) Space Telescope Imaging Spectrograph (STIS) in the 1320 to 1510 \AA\/ region. By cross-correlating these spectra we obtained the orbital velocity amplitude of the companion with an uncertainty in the companion amplitude of 1.9 km sec$^{-1}$. This provides a mass ratio of the Cepheid to the companion of 2.1. The ultraviolet energy distribution of the companion provides the mass of the companion, yielding a Cepheid mass of 5.2 $\pm$ 0.3 M$_\odot$. This mass requires some combination of moderate main sequence core convective overshoot and rotation to match evolutionary tracks. | Masses are the most fundamental parameter governing the evolution of single stars. Interactions between stars in binary/multiple systems can, of course, alter a mass in interesting ways. The tension between the masses derived from evolutionary calculations and those from pulsation calculations has been designated ``the Cepheid mass problem''. A good summary of the differences and the recent state is provided in Neilson, et al (2011), who conclude that it still exists at the 10-20\% level. This implies uncertainty in the evolutionary predictions of luminosity for post-main sequence He burning stars. For classical Cepheids evolutionary calculations are also important in making any adjustments needed to the Leavitt (Period-Luminosity) Law for differences in metallicity between galaxies. Observed masses are needed to clarify these questions. In the Milky Way (MW) there are no Cepheids known in eclipsing binaries. The advent high resolution spectra in the ultraviolet (UV) from satellite observations (originally the International Ultraviolet Explorer [IUE] and currently the Hubble Space Telescope [HST]) has provided orbital velocity amplitudes of the hot companions of several Cepheids. Combining this amplitude with the ground-based orbital velocity amplitude for the Cepheid, and a mass of the companion from the energy distribution in the ultraviolet provides a Cepheid mass. In addition, a dynamical mass has been determined for Polaris using HST astrometry (Evans, et al. 2008, 2018). An upper limit to the mass for W Sgr has been derived from HST astrometry. A summary of the references and results is provided by Evans, et al. (2011). In several cases improved masses are anticipated soon, largely because of the incorporation of interferometry to provide additional resolved orbits. The first result of this program is V1334 Cyg (Gallenne, et al. 2018). Not only is the determination of masses in the MW improving, an additional valuable comparison has become possible with Cepheids in the Large Magellanic Cloud (LMC). Several eclipsing binaries have been discovered in the LMC (Gieren, et al. 2015; Pilecki, et al. 2013; Pilecki, et al. 2015; summarized by Pilecki et al 2018). Thus a comparison of the mass luminosity relation can be made at two metallicities. The first step in mass determination is the derivation of a binary orbit for the primary (Cepheid) from ground-based spectra, which is available for many stars. An early result from UV studies of the companions is that a substantial fraction of the companions are themselves binaries (e.g. Evans, et al. 2005). This is to be expected in high and intermediate mass systems, but the additional observations needed to determine a mass are often prohibitively expensive of telescope time. The system containing the Cepheid V350 Sgr = HD 173297 is one where previous UV observations found the companion to be single. It was observed twice with the HST Goddard High Resolution Spectrograph (GHRS) medium resolution (R = $\lambda$/$\Delta\lambda$ $\approx$ 20,000) between 1840 and 1880 \AA\/ in 1995 (Evans, et al. 1997). From the velocity difference between phases of these two spectra, and the velocity difference from the Cepheid orbit, they derived a mass ratio $M_{Cep}/M_{Comp}$ = 2.1 $\pm$ 0.3. Using the mass from the UV energy distribution of the companion (B9.0 V; Evans and Sugars 1997), they derived a mass for the Cepheid of 5.2 $\pm$ 0.9 M$_\odot$. Since that discussion, a number of factors have contributed to an improved analysis of the system. A new orbit has been derived based on considerable additional velocity data (Evans, et al. 2011), particularly including data near minimum velocity. Because the orbital period is very close to 4 years, uniform phase coverage has been difficult to obtain. In the project here HST spectra obtained with the Space Telescope Imaging Spectrograph (STIS) are combined with this new orbit, providing improved velocities of the companion as discussed in the next section. Successive sections below discuss the observations, the details of velocity measurement from these spectra, the companion, and the results and implications of the new measurements. {\bf Gaia will, of course, ultimately be important in improving the distance and mass. However, the current DR2 release does not include binary motion in the solution. To illustrate, the expected parallax based on the distance from the Benedict et al. (2007) Leavitt law is 1.12 mas. The semi-major axis of the orbit (Evans, et al. 2011) a sin i is 1.32 AU, which is 1.48 mas at this distance. Hence the Gaia solution including orbital motion is clearly needed. Because of this, and also concerns about the effect of Cepheid light variation and possibly the brightness of the system, the appropriate solution will come with later Gaia releases. } | Fig~\ref{ml} puts the mass of V350 Sgr in the context of the measured Cepheid masses, and also of theoretical predictions. New masses are available for V350 Sgr (this paper), V1334 Cyg (Gallenne, et al. 2018) and Polaris (Evans, et al. 2018). Other masses for MW Cepheids are from the sources listed in Evans, et al. (2011). (Note that the mass for SU Cyg is a lower limit, and that of W Sgr is an upper limit.) W Sgr and FF Aql incorporate HST Fine Guidance System (FGS) astrometry (Benedict, et al 2007; Evans, et al. 2009). This is a ``before'' picture, since the accuracy of the masses of S Mus and SU Cyg will be improved in near future including the results from interferometry (which will produce an ``after'' picture). The mass for Polaris is preliminary, and will ultimately be improved using CHARA interferometry, but because of the long period of its orbit, this will not be for several years. Luminosities in Fig~\ref{ml} for the MW Cepheids are derived from the Leavitt Law (Period-Luminosity) of Benedict, et al. (2007). Alternately for V350 Sgr, a radius was derived using the modified Balona technique (Rastorguev and Dambis 2011) after carefully correcting for the effect of the companion on the light curve. The resulting luminosity is slightly smaller than that in Fig~\ref{ml}. The LMC Cepheids in eclipsing binaries have recently been rediscussed by Pilecki, et al (2018) and their masses and luminosities are shown in Fig~\ref{ml} . This includes the interesting case of LMC-CEP-1812 which is crossing the instability strip for the first time (the least luminous Cepheid in Fig~\ref{ml}) and may be a merger product (Neilson et al. 2015). It occurs approximately 0.2 in log(L/L$_\odot$) lower than the relation in Fig~\ref{ml} for second and third crossing stars as expected the predictions of Bono et al. (2016) from the comparison of luminosities between crossings. In addition the system LMC-CEP-1718 A and B contains a pair of first overtone pulsators (the two most massive LMC stars in Fig~\ref{ml}). The combination is unusual in that the more massive is less luminous. However this may be explained by the uncertainty in the luminosities. The range of theoretical predictions from evolutionary tracks is also shown. The left (short dash) line is for the metallicity of the LMC; others are for MW metallicity. The right hand line (long dash) shows the prediction for stars with no core convective overshoot on the main sequence (Bono, et al. 2016). As is well known, these predictions produce the lowest luminosity for a given mass. The two lines in the middle illustrate combinations of parameters which can increase the luminosity for a given mass by increasing the size of the central He core after core hydrogen burning. The solid line has moderate convective overshoot added (d$_{over}$ = 0.2 H$_p$ where H$_p$ is the pressure scale height). The dotted line shows recent Geneva calculations (Anderson, et al. 2014) which include both a smaller amount of overshoot (d$_{over}$ = 0.1 H$_p$) and rotation. The value of 0.5 $\omega_{crit}$ (critical velocity) actually represents the effects of a wide range of rotations well. All the predictions in Fig~\ref{ml} are for combined second and third crossings of the instability strip. The improved accuracy of the mass of V350 Sgr confirms that evolutionary tracks without rotation or overshoot predict too low a luminosity for the mass, which is in agreement with other masses in Fig~\ref{ml}. As improved masses become available, other parameters influencing the luminosity in the Cepheid stage will be more tightly constrained. \noindent {\em Acknowledgments}: {\bf The referee's report improved the clarity of presentation.} Support was provided by HST Grant GO-13368.01-A and Chandra X-ray Center NASA Contract NAS8-03060 (to NRE) and HST-GO-13368.008-A (to CP) . A. Rastorguev acknowledges Russian Foundation for Basic Research (RFBR grant 18-02-00890) for partial support. | 18 | 8 | 1808.10472 |
1808 | 1808.07504_arXiv.txt | The InfraRed Imaging Spectrograph (IRIS) is the first-light client instrument for the Narrow Field Infrared Adaptive Optics System (NFIRAOS) on the Thirty Meter Telescope (TMT). IRIS includes three natural guide star (NGS) On-Instrument Wavefront Sensors (OIWFS) to measure tip/tilt and focus errors in the instrument focal plane. NFIRAOS also has an internal natural guide star wavefront sensor, and IRIS and NFIRAOS must precisely coordinate the motions of their wavefront sensor positioners to track the locations of NGSs while the telescope is dithering (offsetting the telescope to cover more area), to avoid a costly re-acquisition time penalty. First, we present an overview of the sequencing strategy for all of the involved subsystems. We then predict the motion of the telescope during dithers based on finite-element models provided by TMT, and finally analyze latency and jitter issues affecting the propagation of position demands from the telescope control system to individual motor controllers. | \label{sec:intro} % The InfraRed Imaging Spectrograph (IRIS) will be the first-light workhorse instrument for the Thirty Meter Telescope (TMT). IRIS is fed by the Narrow Field Infrared Adaptive Optics System (NFIRAOS), which will provide diffraction-limited optical performance in J, H, and K bands, across the full 35\,arcsec$\times$35\,arcsec field-of-view (FOV) of the IRIS imager. Pickoff optics near the center of the IRIS FOV feed an integral field spectrograph (IFS), which can be run in parallel with the imager. NFIRAOS is a Laser Guide Star (LGS), Multi-conjugate Adaptive Optics System (MCAO) which, in its primary observing mode, uses an asterism of six artificial sodium laser guide stars [observed with corresponding Laser Guide Star Wavefront Sensors (LGS WFS)] to drive two deformable mirrors (DMs), conjugated to 0\,km (DM0) and 11.8\,km (DM11), at a rate of up to 800\,Hz. An additional natural guide star wavefront sensor (NGS WFS) within NFIRAOS provides reference ``truth'' measurements to compensate errors in the LGS system on longer timescales (e.g., radial Zernike error introduced by variations in the structure of the sodium layer). Light is steered on to the NGS WFS using a Star Selection Mechanism (SSM) which combines a movable X-Y stage, and a separate tip/tilt/focus stage. Low-order wavefront error measurements are provided by positionable natural guide star (NGS) On-Instrument Wavefront Sensors (OIWFS) and On-Detector Guide Windows (ODGW) within IRIS. These OIWFSs and ODGWs will be used to measure tip, tilt, and focus errors to which the LGS WFS is blind. NFIRAOS and IRIS are required to deliver 50\% sky coverage at the North Galactic Pole. In this context sky coverage refers to the ability of the OIWFS and ODGW to provide sufficient information to control the low-order modes and meet the TMT LGS MCAO wavefront error (WFE) budget. The OIWFS and ODGW will also control the six plate scale modes at the science focal plane. In addition to the DMs, NFIRAOS can use several slower mechanisms to compensate for longer timescale, and larger amplitude errors. DM0 is mounted to a tip/tilt stage (TTS) with a control bandwidth of approximately 20\,Hz, and acts as a ``woofer'' to compensate slow and large-amplitude tip/tilt errors, while the DM itself acts as a ``tweeter'' for the higher-frequency and lower-amplitude errors. Longer timescale pointing and tip/tilt errors are offloaded to the TMT Telescope Control System (TCS) for compensation. \begin{figure}[ht] \begin{center} \includegraphics[width=\linewidth]{hierarchy.pdf} \end{center} \caption{\label{fig:hierarchy} A representative portion of the observatory software hierarchy relevant to AO operations using IRIS and NFIRAOS. From top to bottom, the layers of this hierarchy are referred to as: the Monitoring and Control Layer; the Sequencing Layer; the Assembly Layer; the Hardware Control Layer; and finally the physical Hardware Layer. Subsystems are color-coded: red for AO; blue for IRIS; and purple for the TCS. Black lines indicate the configuration hierarchy (e.g., sequencers configure assemblies, and assemblies configure HCDs). Note that while the OIWFS/ODGW assemblies are part of IRIS, they are primarily configured by the AOSq. Event Service communication is shown with purple dashed lines; for example, the PK publishes demands for various assemblies within NFIRAOS, IRIS, and the telescope mount, and the RTC publishes corrections for many of the same positioners. A separate high-speed network and protocol, depicted with red dashed lines, is used by the RTC to receive data directly from the WFS, send commands to wavefront correctors, and exchange data with the RPG. The components highlighted in bold are included in the prototype described later in this paper.} \end{figure} Observing with NFIRAOS and IRIS is a complex task, involving communication and feedback between a number of telescope subsystems at different rates, as shown in Fig.~\ref{fig:hierarchy}. Within NFIRAOS is a Real Time Controller (RTC) that is responsible for processing measurements from all of the wavefront sensors (WFS), and providing demands to all of the above mentioned wavefront correctors up to the peak 800\,Hz rate. Statistics collected by the RTC are sent to the Reconstructor Parameter Generator (RPG) which periodically revises and uploads ``control matrices'' to the RTC, which it in turn uses as part of its real-time wavefront corrector calculations. The communication between WFSs, wavefront correctors, the RTC and the RPG use a dedicated high-bandwidth, low-latency network and protocol, and highly-optimized software written in C to meet the stringent real-time requirements. The NFIRAOS RTC is discussed elsewhere in these proceedings\cite{dunn2018}. In contrast, the NGS WFS, OIWFSs, and ODGWs are positioned using continuous demand streams that are sent by the TCS, which converts star coordinates into focal plane coordinates using its pointing kernel (PK), and with error correction streams optionally provided by the RTC. The TCS also provides position streams for numerous other devices throughout the observatory, including the telescope mount. These streams are typically provided at rates 20--100\,Hz, depending on the requirements of the device. These demands are sent using the TMT ``Event Service''. Overall coordination and configuration of the various subsystems at TMT will be orchestrated by ``sequencers'', as depicted in Fig.~\ref{fig:hierarchy}. NFIRAOS and the RTC, as part of the TMT Adaptive Optics (AO) system, will be configured by the Adaptive Optics Sequencer (AOSq). The TCS and science instruments will also have their own sequencers. The sequencers convert requests from higher-level subsystems (e.g., ``configure the AO and science instruments for a particular observation'', ``slew to the location of the observation'', ``engage the AO system'', and finally ``commence science exposures'') into sequences of commands relayed by the TMT Command Service to particular components within each subsystem for which the sequencer is responsible (e.g., configure individual WFS detectors, slew the telescope mount, move filter wheels). The software components with which sequencers interact are called assemblies, with each assembly typically exposing a simple unified interface for a single device (e.g., a detector, filter wheel, beam splitter etc.). Finally, assemblies interact with Hardware Control Daemons (HCD) that directly interface with physical hardware. For example, an HCD for a motor controller may accept commands from an assembly to drive a stage a particular number of millimeters, and then convert the request into motor counts and forward the request to the controller using its particular communications protocol. All of the TMT software systems mentioned above are being developed in parallel, but with staggered schedules. TMT Common Software\cite{gillies2016} (CSW), which includes the Event and Command Services, as well as a framework for building software components such as assemblies and HCDs, passed its final design review in Jan. 2017, and a production version is currently being developed by an India-based vendor named ThoughtWorks, under contract to the India TMT Co-ordination Centre (ITCC). Software for NFIRAOS (being developed at NRC Herzberg), including the RTC, and slow opto-mechanical mechanisms, such as the tip/tilt stage, and NGS WFS mentioned above, recently passed their Final Design Reviews. AO Executive Software (AOESW), is being developed at TMT; the AOSq is in the Preliminary Design Phase, and the RPG the Final Design Phase. IRIS software (led by NRC Herzberg, with U.S. and Japanese partners) has just commenced its Final Design Phase. The TCS [developed at TMT, primarily by the ITCC with support of the Inter-University Centre for Astronomy and Astrophysics (IUCAA)] is currently in its Preliminary Design Phase. Finally, the remaining observatory software work packages will commence their Preliminary Design Phases this autumn. Early prototyping efforts that involve multiple subsystems are therefore critical to gain confidence in the overall observatory software architecture. To this end, this paper describes motion analysis and simulations that have informed the designs of NFIRAOS and IRIS, making use of simulations and prototype software provided by TMT. In particular, we examine one of the more challenging observational scenarios that will be faced by the slow opto-mechanical mechanisms within NFIRAOS and IRIS: closed-loop dithering. In order to cover large areas of sky efficiently, TMT will sequentially offset up to 30 arcsec on-sky, performing separate science integrations at each location, as part of a single observing sequence. The NGS WFS and OIWFS positioners must undergo complementary motion to that of the telescope during these offsets so that their guide stars can be observed continuously. If the WFSs lose lock on their guide stars, NFIRAOS will cease performing AO corrections, and a substantial AO re-acquisition time penalty will be incurred upon completion of the dither offset (potentially as much as tens of seconds). Such a penalty would render dithering (and therefore large-area mapping) incredibly inefficient. In Section~\ref{sec:motion} the expected motion of the telescope during a dither is analyzed. The results of this analysis are used to justify a TCS demand publication rate 20\,Hz for the control of NFIRAOS and IRIS WFS positioners. Then, Section~\ref{sec:swproto} uses CSW framework and service prototypes [which run on Java Virtual Machines (JVM)], to create a ``vertical slice'' through the control system, including the TCS, an IRIS assembly, an HCD, and a real motion controller. Timing measurements are used to explore latency and jitter in the control system, ultimately demonstrating that the motion requirements can be met. | In this paper we have demonstrated the feasibility of closed-loop AO dithering with IRIS and NFIRAOS from the perspective of their slow opto-mechanical mechanisms. Taking an approach in which the telescope mount and wavefront sensor positioners within IRIS and NFIRAOS are driven with a smooth ``S'' profile produced by the TCS at 20\,Hz, the expected position uncertainties are within 0.2\,mm in the focal plane. This limit was established to avoid saturating the NFIRAOS tip/tilt stage (which compensates for pointing errors in real-time). Key elements of the control system for one of the positioners were prototyped in Java, including the propagation of demands using Common Software Services supplied by TMT, and interaction with a real Galil 8-axis Motion Controller. An essential component of our control strategy is linear extrapolation of past TCS demands by the low-level Hardware Control Daemon (HCD) evaluated at the precise instants they are required. Finally, we study jitter in the system, finding that significant ($\gtrsim$20\,ms) pauses can generally be associated with activities of the Java garbage collector (GC). Switching from the default settings used in Java 8 (Parallel garbage collector), to the newer Garbage First Garbage Collector (G1GC), with the inclusion of the \texttt{MaxGCPauseMillis} parameter set to 10\,ms, removes essentially all of the significant pauses measured over 48\,hr. While further tuning may continue to improve the jitter performance, we note that the greatest outliers in our Java test applications only slightly exceed those of a reference application written in C that does not use a GC. A number of future tests are planned. First, the HCD may be updated to either: (a) drive a real motor; or (b) include an improved simulation, with a real PID loop, and at least trapezoidal velocity profiles (as the current simulation allows instantaneous velocity changes). The software components and CSW services should be run on separate physical servers to include the effects of the network. Real-time performance can be improved by setting real-time priorities, core isolation, and the use of NUMA regions; in particular, it would be interesting to see if the width of the reference C application control loop period histogram can be reduced. If it is possible to significantly improve the C application performance, it may then be worth further experimentation with the many GC tuning parameters. Finally, once the production TMT CSW services are released, these test applications should be re-written to make use of them, and re-benchmarked. | 18 | 8 | 1808.07504 |
1808 | 1808.01819_arXiv.txt | We demonstrate that the existing neutron-star cooling data can be appropriately described within ``the nuclear medium cooling scenario" including hyperons under the assumption that different sources have different masses. We use a stiff equation of state of the relativistic mean-field model MKVORH$\phi$ with hadron effective couplings and masses dependent on the scalar field. It fulfills a large number of experimental constraints on the equation of state of the nuclear matter including the $2\,M_{\odot}$ lower bound for the maximum predicted neutron-star mass and the constraint for the pressure from the heavy-ion particle flow. We select appropriate $^1S_0$ proton and $\Lambda$ hyperon pairing gap profiles from those exploited in the literature and allow for a variation of the effective pion gap controlling the efficiency of the medium modified Urca process. The $^3P_2$ neutron pairing gap is assumed to be negligibly small in our scenario. The possibility of the pion, kaon and charged $\rho$-meson condensations is for simplicity suppressed. The resulting cooling curves prove to be sensitive to the value and the density dependence of the pp pairing gap and rather insensitive to the values of the $^1S_0$ neutron pairing gaps. | During many years the problem of the cooling of neutron stars (NSs) attracted great interest, e.g. see \cite{Bahcall:1965zzb,Tsuruta:1979,ST83,Migdal:1990vm,Yakovlev:2000jp,Voskresensky:2001fd,Page:2004fy, Potekhin:2015qsa,Schmitt:2017efp,Sedrakian:2018ydt}. After first tens of seconds, or at most several hours for most massive NSs, the typical temperature of a newly born NS decreases below the neutrino-opacity temperature $T_{\rm opac}\sim $~1~MeV \cite{ST83,Sawyer:1978qe,Voskresensky:1986af,Migdal:1990vm,Voskresensky:2001fd}. Below we will be interested in the stage when $T<T_{\rm opac}$. At this stage that lasts for the first $\sim 10^5$ yr. NSs are cooled by the direct neutrino radiation from their interiors and then, for $t\gsim 10^6$ yr., by the photon radiation from the surface. For $T<T_{\rm opac}$ the typical averaged neutrino energy is $\sim$ several $T$, which is much larger than the nucleon particle width $\Gamma_N \sim T^2/\varepsilon_{{\rm F}}$. Therefore nucleons can be treated within the quasiparticle approximation~\cite{Knoll:1995gs,Knoll:1995nz,Voskresensky:2001fd}, and the efficiency of processes of neutrino production can be graded simply according to their available phase space. The direct Urca (DU) one-particle reactions, e.g. $n\to pe\bar{\nu}$, yield the largest neutrino emissivity, $\epsilon_{\nu}^{\rm DU}\sim 10^{27} T_9^6\Theta(n-n_{c,N}^{\rm DU})$ $\frac{{\rm erg}}{{\rm s}\cdot {\rm cm}^3}$ for non-superfluid systems, $T_9 =T/10^9$ K, $\Theta(x)$ is the step-function, cf. \cite{Lattimer:1991ib}. Thereby the DU processes on nucleons exist only in NSs with masses $M>M_{c}^{\rm DU}$ in which the central density exceeds the value $n_{c,N}^{\rm DU}$. An estimate of the DU-threshold follows from the triangle inequality for momentum conservation, that demands the proton concentration in the $npe\mu$ matter to be above 11.1-14.8\%. Therefore the DU processes set in at very different critical densities depending on the density dependence of the symmetry energy, which is an important quantity for the description of both heavy-ion reactions and NSs. The DU emissivity proves to be so high that, if DU processes occurred in majority of NSs, the latter could not be visible in soft $X$-rays. It should, however, not be expected that the objects observed in X-rays are some exotic family of NSs rather than typical NSs. As a "weak DU constraint" for an equation of state (EoS) \cite{Klahn:2006ir} suggested to require that $M_{c}^{\rm DU}>1.35 \, M_{\odot}$ since measured masses of the majority of NSs in binaries lie between $(1.35 -1.4) \, M_{\odot}$. On the other hand, according to the pulsar population modeling most of NSs have masses below $1.5\,M_\odot$, see \cite{Alsing:2017bbc}, Figs. 1, 2 there. Thereby \cite{Kolomeitsev:2004ff,Blaschke:2004vq,Klahn:2006ir} suggested to require the absence of the DU processes in NSs with $M<1.5 \, M_{\odot}$. Reference \cite{Klahn:2006ir} named the requirement $M_{c}^{\rm DU}>1.5 \, M_{\odot}$ a ``strong DU constraint". The DU constraint puts a restriction on the density dependence of the symmetry energy. The modified Urca (MU) two-nucleon processes have a much smaller emissivity than DU. Among them the luminosity of the neutron branch $nn\to npe\bar{\nu}$ of MU process is greater than that of the proton branch $np\to ppe\bar{\nu}$ and the nucleon bremstrahlung (NB) $nn\to nn\nu\bar{\nu}$, $np\to np\nu\bar{\nu}$, cf. \cite{Tsuruta:1979, ST83,Friman:1978zq}. The evaluation of the MU and the NB neutrino emissivities requires to know the nucleon-nucleon ($NN$) interaction amplitude in the baryon medium. For the $NN$ interaction Friman and Maxwell \cite{Friman:1978zq} used the free one-pion exchange (FOPE) model. The density dependence of the reaction rates calculated with the FOPE model is very weak and thereby the neutrino radiation from a NS in this model depends very weakly on the star mass. One gets an estimate of the emissivity of the MU processes, $\epsilon_{\nu}^{\rm MU}\sim 10^{21} T_9^8$ $\frac{{\rm erg}}{{\rm s}\cdot {\rm cm}^3}$. This naive estimate used in many works is essentially modified, if various in-medium effects are taken into account in the $NN$ interaction amplitude \cite{Voskresensky:1985qg,Voskresensky:1986af,Senatorov:1987aa,Voskresensky:1987hm, Migdal:1990vm,Voskresensky:2001fd}. Another important point is that the nucleons may form $nn$ and $pp$ Cooper pairs, cf. \cite{Migdal1959,ST83,Senatorov:1987aa,Voskresensky:1987hm,Migdal:1990vm,Yakovlev:1999sk, Yakovlev:2000jp,Voskresensky:2001fd,Sedrakian:2018ydt}, if the temperature is lower than corresponding critical temperatures $T^{nn}_c$ and $T^{pp}_c$, respectively. Neutrons undergo $^1S_{0}$ pairing for $n\lsim (0.6-0.8) \, n_0$ and $^3P_2$ pairing for $0.8 \, n_0 \lsim n \lsim (3-4) \, n_0$, whereas protons are paired in $^1S_{0}$ state for $n\lsim (2-4) \, n_0$, where $n_0 = 0.16$ fm$^{-3}$ is the nucleon saturation density. Typical values for the $nn$ and $pp$ pairing gaps, $\Delta_{nn}$ and $\Delta_{pp}$, vary in the range $\sim (0.1-{\rm several})$~MeV, cf. \cite{HGRR1970,Tamagaki1970} and the recent review \cite{Sedrakian:2018ydt} for further references. The critical temperature for the $^1S_0$ pairing is $T^{NN}_c \simeq 0.57 \Delta_{NN}$. Due to an exponential dependence on the value of the $NN$ interaction amplitude in the particle-particle channel, existing estimates of the pairing gaps are rather uncertain. The values of the $^3P_2$ neutron pairing gaps are known especially poorly. BCS-based estimates \cite{HGRR1970,Takatsuka:2004zq} yield values similar to those for the $^1S_0$ pairing. However, taking into account a medium-induced spin-orbit interaction leads to a tiny value of $\Delta({^3P_2}) \lsim 10$ KeV~\cite{Schwenk:2003bc}. Within the same scenario as in this work the dependence of the cooling curves on the values of $^3P_2$ neutron and $^1S_0$ proton paring gaps was studied in~\cite{Grigorian:2005fn}. In this work a reasonable fit of the compact star cooling data was obtained for a strongly suppressed value of the $^3P_2$ neutron pairing gap, in favour of results of ~\cite{Schwenk:2003bc}. The emissivities of the reactions with participation of nucleons are suppressed by the so-called $R$-factors, describing the available phase space for a reaction. Typically the emissivity of the DU processes is suppressed by $e^{-(\max\{\Delta_{nn}, \Delta_{pp}\}/T)}$. Nevertheless the DU processes remain relatively rapid even in presence of the nucleon superfluidity, because the emissivities of the MU and NB processes are suppressed even more, e.g. for the neutron branch of the MU process a typical suppression factor is $e^{-(\Delta_{nn}+ \Delta_{pp})/T}$, cf. \cite{ST83,Maxwell:1979zz}. In a realistic calculation the R-factors have essentially more complicated forms, cf. \cite{Yakovlev:1999sk,Yakovlev:2000jp}, which we take into account in our scenario. In the historically first so-called {\em{standard scenario}} of NS cooling \cite{Tsuruta:1979,ST83} the processes were calculated without taking in-medium effects into account. By that time only upper limits on NS surface temperatures were measured by the Einstein observatory at an assumption that there exist NSs in the observed supernova remnants. Some of these upper limits proved to be high, other low but it was not known whether the NSs exist in those remnants. For the ideal Fermi gas of nucleons \cite{Bahcall:1965zzb} the DU processes are not allowed by the energy-momentum conservation. Thereby the MU processes were considered as the most important channel for relevant values of internal temperatures, $T\sim 10^{8}$--$10^{9}$~K for $t\lsim 10^5$ yr. The MU and NB processes were evaluated using the FOPE model of Friman and Maxwell \cite{Friman:1978zq}. With the MU and NB processes it was possible to explain the highest of measured upper limits on the surface temperatures of NSs, see Fig. 11.3 in \cite{ST83}. To explain the lowest measured upper limits at assumption of existence of NSs in those remnants one exploited a possibility of the $P$-wave pion condensation in NS interiors suggested by A.B. Migdal, see \cite{Migdal:1978az}. In presence of a pion condensate pion Urca (PU) processes $N_1\pi_c \to N_2 e\bar{\nu}$ become possible for $n>n_c^{\rm PU}> n_0$. Their emissivity is roughly estimated as $\epsilon_{\nu}^{\rm PU}\sim 10^{26} T_9^6\Theta (n-n_c^{\rm PU})$ $\frac{{\rm erg}}{{\rm s}\cdot {\rm cm}^3}$ for non-superfluid systems, cf. \cite{Maxwell:1977zz}. Like the DU processes, the PU processes are of the one-nucleon origin, but the requirement of energy-momentum conservation can be always fulfilled for $n>n_c^{\rm PU}$ by absorbing momentum by the condensate. In the eighties of the previous century the prevailing opinion was that all NS masses should be close to the values for known binary radio pulsars, $(1.35 - 1.4)\, M_{\odot}$, cf. \cite{ST83}. Under this assumption within the standard scenario two possibilities were considered. First one was that in supernova remnants for which the low upper limits on the NS surface temperatures were put there are no NSs, and that there is no pion condensation in NSs. Second one was that in at least some supernova remnants for which the low upper limits on the NS surface temperatures were put there are NSs, and thereby there is pion condensation in NSs. The existed data did not allow to determine whether the observable objects are the NSs, slowly cooled mainly by the MU processes, or the rapid coolers with the PU processes enabled. Besides the DU, MU and MN reactions a nucleon pair-breaking-formation (PBF) reaction channels $N\to N\nu\bar{\nu}$, $N=n,p$, are opened up for $T<\{T^{nn}_c, T^{pp}_c\}$ respectively \cite{Flowers:1976ux,Senatorov:1987aa, Voskresensky:1987hm}. The nPBF neutrino process in presence of the $nn$ pairing was introduced first in \cite{Flowers:1976ux} to occur on the vector current. Its emissivity was estimated without inclusion of in-medium effects in vertices of the process to be of the order or smaller than that for the MU processes. References \cite{Senatorov:1987aa, Voskresensky:1987hm} performed calculations of the neutrino emissivity in the non-equilibrium Green function technique that naturally allows to separate various processes within the quasiparticle approximation valid for $T\ll E_{{\rm F},N}$, where $E_{{\rm F},N}$ is the nucleon Fermi energy. The authors considered emissivity of the nPBF process and suggested a possibility of the pPBF one in presence of the $pp$ pairing. They found that the emissivity of the PBF processes is roughly estimated as $\epsilon_{\nu}^{\rm PBF}\sim 10^{28} [\frac{\Delta_{NN}}{{\rm MeV}}]^7 \sqrt{T/\Delta_{NN}} e^{-(2\Delta_{NN}/T)}$ $\frac{{\rm erg}}{{\rm s}\cdot {\rm cm}^3}$ that exceeds the emissivity of the MU processes for $\Delta_{NN}\gsim 10^9$ K. The effect of the PBF reactions on the cooling was first incorporated in the cooling code in \cite{Schaab:1996gd}. Then these processes were included in all existing cooling codes. Already in \cite{Senatorov:1987aa, Voskresensky:1987hm} an important role of in-medium effects in PBF processes was pointed out, especially for the pPBF reactions. Further, Leinson and Perez \cite{Leinson:2006gf} noticed that in the vector current channel incorporation of in-medium effects needs a special care due to a necessity to fulfill the Ward-Takahashi identities (purely in-medium effect!). Thereby the emissivity of the nPBF process on the vector current in presence of the $^1S_0$ $nn$ pairing proves to be dramatically suppressed, by a pre-factor $v_{\rm F}^4$, where $v_{\rm F}$ is the nucleon Fermi velocity. Detailed analyses of the PBF reactions \cite{Kolomeitsev:2008mc,Kolomeitsev:2010hr,Kolomeitsev:2011wz} have shown that by taking into account the in-medium dressing of vertices, the main contribution to the PBF emissivity in presence of the $^1S_0$ $nn$ and/or $pp$ pairings comes from processes on the axial current, and the emissivity is thereby suppressed only as $v_{\rm F}^2$ rather than as $v_{\rm F}^4$. Also, \cite{Leinson:2011jr} noticed that in the presence of the $^3P_2$ $nn$ pairing the nPBF emissivity might be not suppressed by the $v_{\rm F}^2$ factor, in favor of numerical estimates \cite{Voskresensky:1987hm} previously used in the neutron-star cooling simulations, cf. \cite{Schaab:1996gd}. Recall here that, if the $^3P_2$ $nn$ pairing gap is $\lsim 10$ KeV as follows from estimates ~\cite{Schwenk:2003bc}, the nPBF processes are not effective at all for $t\lsim (10^5 -10^6)$ yr of our interest. At present time there exists information on surface temperature-time ($T_s -t$) dependence for many pulsars. The $T_s-t$ data can be separated in three groups dubbed as slow cooling, intermediate cooling and rapid cooling. Recently the Chandra observatory measured the suface temperature of the young compact star in the remnant of the historical supernova Cassiopeia A (Cas A) of the year 1680, cf. \cite{Tananbaum:1999kx,Hughes:1999ph}. The description of the cooling of this object caused a lovely discussion \cite{Page:2010aw,Shternin:2010qi,Blaschke:2011gc,Blaschke:2013vma, Elshamouty:2013nfa,Ho:2014pta,Grigorian:2016leu}. Besides, the cooling model must also explain the hot central compact object in the supernova remnant XMMU J173203.3-344518, cf. \cite{Klochkov}. Note that in order to explain the difference in the cooling of the slowly and rapidly cooling objects a three order of magnitude difference in their luminosities is required \cite{Voskresensky:2001fd}. Basing on the standard scenario the so-called {\em{minimal cooling paradigm}} was formulated \cite{Page:2004fy,Page:2009fu}. As in the standard scenario, in this approach the medium effects are assumed to play only a minor role and masses of NSs are assumed to be close to $1.4M_{\odot}$. In this scenario the difference in the measured values of the surface temperatures of various sources of soft $X$ rays is supposed to be explained by the heterogeneity in envelope compositions for the young stars: light element compositions for some of them and heavy element compositions for others. A similar approach has been extensively used by other researches, cf. \cite{Yakovlev:1999sk,Yakovlev:2000jp,Tsuruta:2002ey,Page:2010aw,Shternin:2010qi,Elshamouty:2013nfa,Ho:2014pta,Potekhin:2015qsa,Raduta:2017wpp} and refs. therein. The MU and NB processes in these works were considered within the FOPE model of \cite{Friman:1978zq}. The nPBF processes for the $^1S_0$ pairing on the vector current were included first, as in \cite{Flowers:1976ux}, without in-medium modification factors and later the processes on the vector current were suppressed following \cite{Leinson:2006gf} and the processes on axial current were included with taking into account the $v_{\rm F}^2$ suppression factor as in \cite{Kolomeitsev:2008mc,Kolomeitsev:2010hr}. The emissivity of the pPBF processes calculated without in-medium modification is very small, so these processes are considered as not effective in the minimal cooling scenario. In reality due to in-medium modification of the vertices these processes might be almost as effective as nPBF ones \cite{Kolomeitsev:2008mc,Kolomeitsev:2010hr}, depending on values of $nn$ and $pp$ pairing gaps. The values of the $nn$ gap for the $^3P_2$ $nn$ pairing and the $^1S_0$ $pp$ gaps were found to better fit the data on the surface temperatures of the pulsars. For the $^3P_2$ $nn$ pairing values $\Delta_{nn}\sim 10^9$~K were suggested. The DU processes were assumed not to occur. Thereby the main cooling agents in this scenario are the MU, NB and nPBF processes the latter going on the neutrons paired in the $^3P_2$ state. It proves to be that the existing data on the time dependence of the surface temperatures of pulsars are hardly explained within the minimal cooling scenario. In particular, the hot object XMMU J173203.3-344518 is not explained. The agreement with other data can be achieved only if the scenario is supplemented by a possibility of the efficient DU reaction in majority of NSs. Under the outdated assumption that masses of all the NSs with measured surface temperatures are close to $1.4 \, M_{\odot}$, the DU process should be effective already for $M\simeq 1.4 \, M_{\odot}$. However the latter assumption disagrees with a broad distribution of NSs over the masses, as follows from the population synthesis modeling and from supernova simulations \cite{Popov:2004ey,Alsing:2017bbc}. Moreover, recent measurements of masses of the heaviest binary pulsars demonstrated that the maximum compact star mass should be $>2 \, M_{\odot}$. It was found that PSR J1614-2230 has the mass $M = 1.928 \pm 0.017 \, M_\odot$ \cite{Demorest:2010bx,Fonseca:2016tux} and PSR~J0348+0432, the mass $M = 2.01 \pm 0.04 \, M_\odot$~\cite{Antoniadis:2013pzd}. The measurements of the high masses of the pulsars PSR J1614-2230 \cite{Demorest:2010bx,Fonseca:2016tux} and PSR J0348-0432 \cite{Antoniadis:2013pzd} and of the low masses for PSR J0737-3039B \cite{Kramer:2006nb} for the companion of PSR J1756-2251 \cite{Faulkner:2004ha,Ferdman} and the companion of PSR J0453+1559 have provided the proof for the existence of NSs with masses varying at least from $1.2$ to $1.97 \, M_{\odot}$. Thus, to explain existence of heaviest NSs the EoS of the NS matter should be sufficiently stiff, cf. \cite{Klahn:2006ir,Voskresenskaya:2012np,Gandolfi:2015jma,Lattimer:2015nhk}. Already long ago Refs. \cite{Voskresensky:1985qg,Voskresensky:1986af} suggested that different NSs have different masses and the $T_s -t$ history of NSs of various masses should be essentially different. With this assumption first the Einstein observatory data were explained and then it was shown that also Chandra observations can be naturally explained. Refs. ~\cite{Voskresensky:1985qg,Voskresensky:1986af,Senatorov:1987aa, Voskresensky:1987hm,Migdal:1990vm,Voskresensky:2001fd} suggested that for $n\gsim n_0$ the $NN$ interaction amplitude is mainly controlled by the soft pion exchange. The pion softening effect, which increases with increase of the density, is the consequence of the strong polarization of the dense nuclear matter. The authors calculated the rates of the two-nucleon processes with inclusion of in-medium effects and found a strong dependence of the $NN$ interaction amplitude on the density (and respectively mass of the object). Thus the FOPE-based MU diagram \begin{center} {\includegraphics[width=.2\textwidth]{YaF-borndiag.pdf}} \label{pairneutr} \end{center} (the zig-zag line corresponds to the free pion and the dots denote free vertices) should be replaced by the medium one-pion exchange (MOPE) diagrams of the medium modified Urca (MMU) processes: \begin{center} {\includegraphics[width=.2\textwidth]{YaF-MMU3.pdf}}\,,\quad \quad {\includegraphics[width=.2\textwidth]{YaF-MMU1.pdf}},\,\,\quad\quad {\includegraphics[width=.2\textwidth]{YaF-MMU2.pdf}},\,\, \label{MMU123} \end{center} where bold-wavy line corresponds to the dressed pion and the hatched vertex takes into account the $NN$ correlations. The first diagram naturally generalizes the MU (FOPE) contribution (\ref{pairneutr}). Besides the MOPE, the $NN$ interaction amplitude contains also a more local interaction part. However the correlation effects on the local $NN$ interaction lead to a suppression of the amplitude. Therefore for $n\gsim n_0$ the main contribution is given by the MOPE. As the result of the pion softening, the $NN$ interaction amplitude proves to be strongly enhanced for $n>n_0$, \cite{Voskresensky:1986af,Migdal:1990vm,Voskresensky:2001fd}, whereas for $n\ll n_0$ the same polarization effects result in a suppression of the $NN$ amplitude compared to that computed in the FOPE model, cf. \cite{Blaschke:1995va,Hanhart:2000ae,Knoll:1995nz}. Evaluations have shown that first diagram (\ref{MMU123}) gives a smaller contribution to the MMU emissivity for $n \gsim n_0$ than the second and third diagrams, which incorporate processes occurring through intermediate reaction states. Note that the latter two diagrams do not contribute, if one approximates the nucleon-nucleon interaction by a two-body potential. Also the second and third diagrams (\ref{MMU123}) do not contribute to the emissivity of the medium-modified nucleon bremstrahlung (MNB) processes due to symmetry arguments. All these peculiarities of the density dependence of the $NN$ interaction amplitude were taken into account in the mentioned works, see \cite{Voskresensky:1986af,Migdal:1990vm,Voskresensky:2001fd}. For $n=n_c^{\rm PU}$, the pion condensation may appear by a first-order phase transition, see \cite{Migdal:1990vm}. Estimates give $n_0<n_c^{\rm PU}\lsim 3n_0$. Evaluations ~\cite{Voskresensky:1986af,Senatorov:1987aa, Voskresensky:1987hm,Migdal:1990vm,Voskresensky:2001fd} showed that at $n=n_c^{\rm PU}$ the MMU emissivity may exceed the MU one calculated with the FOPE model by $3-4$~orders of magnitude. Besides the MMU and MNB processes \cite{Voskresensky:1985qg,Voskresensky:1986af}, the PBF processes \cite{Senatorov:1987aa, Voskresensky:1987hm,Kolomeitsev:2008mc,Kolomeitsev:2010hr,Kolomeitsev:2011wz} and other ones were incorporated. Basing on results of these works the {\em{nuclear medium cooling scenario}} was developed. The results of the calculations were confronted to the data on the $T_s-t$ plane \cite{Schaab:1996gd,Blaschke:2004vq,Grigorian:2005fn,Blaschke:2011gc,Blaschke:2013vma, Grigorian:2016leu} demonstrating an overall agreement, without necessity to assume presence of the DU processes. Moreover, the absence of sub-millisecond pulsars might be naturally explained within the nuclear medium cooling scenario \cite{Kolomeitsev:2014epa,Kolomeitsev:2014gfa}. Two possibilities, one which allows for the PU processes and other one not allowing the pion condensation, were considered. In both cases the $T_s-t$ data can be appropriately described. Near the pion condensation threshold there may also exist a region of the fermion condensation that allows for the neutrino processes with the neutrino emissivity $\propto T^5$, cf. \cite{Voskresensky:2000px}. Besides the $P$-wave pion condensation, in the dense matter may occur the $S$-wave or $P$-wave kaon condensations \cite{Kaplan:1987sc,Kolomeitsev:1995xz,Kolomeitsev:2002pg} and the $S$-wave charged $\rho$ condensation \cite{Voskresensky:1997ub,Kolomeitsev:2004ff,Kolomeitsev:2017gli}. Although all these possibilities lead to an enhanced cooling of NSs \cite{Tatsumi:1988up,Fujii:1993cf,Kolomeitsev:2004ff} the resulting emissivities are typically hundred times smaller than that for the DU process. For the sake of simplicity we focus below on the possibility that the pion softening is saturated for $n\gsim n_c^{\rm PU}\sim 3n_0$ and disregard the possibilities of pion, kaon, charged $\rho$-meson and fermion condensations. First simulations done within the nuclear medium cooling scenario \cite{Schaab:1996gd,Blaschke:2004vq,Grigorian:2005fn,Blaschke:2011gc} exploited the HHJ version of the variational Akmal-Pandharipande-Ravenhall (APR) EoS with a fitting parameter $\delta = 0.2$ \cite{Heiselberg:1999mq} yielding $M_{max}\simeq 1.92 \, M_{\odot}$. In \cite{Blaschke:2013vma} we exploited the HDD EoS, being similar to the APR EoS \cite{Akmal:1998cf} up to $4\,n_0$, but stiffer at higher densities than the HHJ EoS, producing the maximum NS mass compatible with the observations of PSR J1614-2230 \cite{Demorest:2010bx,Fonseca:2016tux} and PSR J0348-0432 \cite{Antoniadis:2013pzd}. In \cite{Grigorian:2016leu} we used still stiffer DD2 and DD2vex EoSs satisfying the constraint $M_{max}>2 \, M_{\odot}$, cf. \cite{Typel:2009sy}. The DD2 EoS and its modification DD2vex EoS are the RMF based EoSs where one uses the density dependence of the couplings. However all these EoSs do not include the possibility of filling of the hyperon Fermi seas. Appearance of hyperons leads to a softening of the EoS and reduction of the maximum NS mass. By employing a hyperon-nucleon potential, the maximum mass of NSs with hyperons was computed ~\cite{Djapo:2008au} to be well below $1.4 M_{\odot}$. Within the RMF approach the critical densities for the appearance of first hyperons prove to be rather low, $n_{\rm c}^{H}\sim 3n_0$, cf.~\cite{Glendenning,SchaffnerBielich:2008kb,Weissenborn:2011kb}, if the hyperon coupling constants satisfying the SU(6) symmetry relations are fitted from the hyperon potentials in the nuclear medium at $n= n_0$. The difference between NS masses obtained with and without inclusion of hyperons proves to be so large for reasonable hyperon fractions in the standard RMF approach, that in order to get the maximum NS mass satisfying experimental constraints one has to start with very stiff purely nucleon EoS. The latter assumption hardly agrees with the results of the microscopically-based variational EoS~\cite{Akmal:1998cf} and the EoS calculated with the help of the auxiliary field diffusion Monte Carlo method~\cite{Gandolfi:2009nq}. The problem was called ``the hyperon puzzle" \cite{SchaffnerBielich:2008kb}. The suggested explanations required additional assumptions, see discussion in~\cite{Fortin:2014mya}, for example, the inclusion of an interaction with a $\phi$-meson mean field, and the usage of smaller hyperon-nucleon coupling constants following the SU(3) symmetry relations~\cite{Weissenborn:2011ut}, as well as other modifications. The above mentioned hyperon puzzle is naturally solved within the RMF EoSs with hadron effective masses and coupling constants dependent on the scalar field $\sigma $ \cite{Maslov:2015msa,Maslov:2015wba}. Appropriate KVORcutH$\phi$ and MKVORH$\phi$ models were constructed. Other important constraints on the EoS, e.g. the flow constraint from heavy-ion collisions \cite{Danielewicz:2002pu,Fuchs,Lynch:2009vc}, are fulfilled within these models as well \cite{Maslov:2015msa,Maslov:2015wba,Kolomeitsev:2016ptu,Kolomeitsev:2017gli}. However, there remains another part of the hyperon puzzle: in presence of hyperons the efficient DU reactions on hyperons, e.g. $\Lambda\to p+e+\bar{\nu}$ become possible, cf. \cite{Maxwell:1986pj}. These reactions are operative only, if the hyperon concentration exceeds some threshold value, that may happen with an increase of the baryon density $n$ above some critical value $n^{\rm DU}_{c, H}$. They accelerate the cooling of NSs with $M>M^{\rm DU}_{c, H} $, where $M^{\rm DU}_{c, H} $ is a NS mass, at which the central density reaches the critical value $n^{\rm DU}_{c, H}$. In the reaction on $\Lambda$ mentioned above the threshold value is only slightly above the density $n_{c,\Lambda}$ of their appearance. Up to now the cooling of NSs with inclusion of hyperons with a stiff EoS satisfying the constraint $M_{max}>2 \, M_{\odot}$ was considered only within the minimal cooling scenario, cf. \cite{Raduta:2017wpp}. As it is seen from the figures of this work, the existing surface temperature-age data are hardly described in the minimal cooling scheme. In the given paper we demonstrate how a satisfactory explanation of existing observational pulsar cooling data is reached within the ``nuclear medium cooling" scenario of \cite{Blaschke:2004vq}, now with the RMF EoS MKVORH$\phi$ with a $\sigma$-scaled hadron effective masses and coupling constants including hyperons \cite{Maslov:2015msa,Maslov:2015wba}. The paper is organized as follows. In the next section we introduce the MKVORH$\phi$ EoS which we then employ in calculation of the cooling history of NSs. Section \ref{sect::cool_input} introduces inputs for the neutrino cooling calculations. Section \ref{sect::results} presents results of numerical calculations. Section \ref{sect::conclusion} contains concluding remarks. The results were briefly announced on the conference ``Compact Stars in the QCD Phase Diagram VI: Cosmic matter in heavy-ion collision laboratories?" (CSQCD VI), cf. \cite{Grigorian:2017xqd}. | \label{sect::conclusion} In the given work we exploited the equation of state of the hadronic MKVOR model (assuming absence of hyperons) and the MKVORH$\phi$ model (with inclusion of hyperons). We have demonstrated that the presently known cooling data can be appropriately described within our nuclear medium cooling scenario both without inclusion of hyperons (MKVOR model) and with included hyperons (within MKVORH$\phi$ model) under assumption that different sources have different masses, provided we use appropriately selected models for proton gaps, e.g. following the models BCLL, CCYms, EEHOr, T, EEHO, CCDK, see left panel of Fig. \ref{Protongaps}, and the models TT1, TTGm, TN-NDSoft, TN-Ehime, TN-FGA for $\Lambda$ pairing gaps, see right panel of Fig. \ref{Protongaps}. To larger densities spreads the proton gap and the higher values it gets within the density interval of its existence, the more regular is the neutron-star mass dependence of the cooling curves in our scenario. The effect of hyperons is rather strong in MKVORH$\phi$ model, after the DU processes with participation of $\Lambda$s are allowed (for $M>1.429\, M_{\odot}$). The so called ``strong DU constraint" of \cite{Klahn:2006ir} that the DU processes should not be allowed in neutron stars with $M< 1.5 \, M_{\odot}$ is not fulfilled for the DU processes with hyperons. However one should bear in mind that the DU emissivity on hyperons is typically two orders of magnitude smaller than that for the DU processes on nucleons. Therefore the strong DU constraint should be partially relaxed. Besides, note that influence of hyperons on the neutron-star cooling would be much less significant, if we used the KVORcut-based models to construct the equation of state, cf. \cite{Maslov:2015wba}. Besides hyperons, another contribution to the neutrino emissivity could come from the processes with $\Delta$-isobars. Among them the density for appearance of $\Delta^-$ is the lowest because it becomes possible to replace the leptons by $\Delta^-$ to fulfill the charge neutrality condition \cite{Kolomeitsev:2016ptu, Li:2018qaw}. However, the threshold baryon density for the DU processes on $\Delta^-$ is larger than the one for the nucleon DU processes, since the $\Delta^-$ fraction is always lower than the proton one \cite{Prakash:1992zng}. Within the MKVOR-based model used in the given work appearance of the $\Delta$s leads to a notable increase of the proton fraction, so the critical density for the nucleon DU processes decreases. Nevertheless the critical value of the neutron-star mass $M^{\rm DU}_{c, p}$ remains higher than 1.5~$M_\odot$, provided the $\Delta$ potential satisfies inequality $U_\Delta (n_0)> -88$~MeV, whereas a realistic value of the $\Delta$ potential is $U_\Delta (n_0)\sim -50$~MeV. On the other hand the hyperon fractions change only a little after inclusion of $\Delta$s \cite{Kolomeitsev:2016ptu}. Thus in this work devoted to the study of the neutron-star cooling in hadronic models with hyperons we disregarded the possibility of appearance of $\Delta$s. In a more detail the latter possibility will be studied elsewhere. We focused on the study of purely hadronic equation of state. Another possibility is to consider cooling of hybrid stars within this equation of state for the hadron phase and allowed the first order phase transition from the hadron to the quark matter, cf. \cite{Blaschke:2000dy,Grigorian:2004jq,Sedrakian:2015qxa}, and permitted various mixed phases, cf. \cite{Voskresensky:2002hu,Maruyama:2005tb,Maruyama:2007ey,Ayriyan:2017nby} and refs therein. We hope to return to such a study in subsequent works. | 18 | 8 | 1808.01819 |
1808 | 1808.08753_arXiv.txt | {One of the important questions of astrochemistry is how complex organic molecules, including potential prebiotic species, are formed in the envelopes around embedded protostars. The abundances of minor isotopologues of a molecule, in particular the D- and $^{13}$C-bearing variants, are sensitive to the densities, temperatures and time-scales characteristic of the environment in which they form, and can therefore provide important constraints on the formation routes and conditions of individual species.} {The aim of this paper is to systematically survey the deuteration and the $^{13}$C content of a variety of oxygen-bearing complex organic molecules on Solar System scales toward the ``B component'' of the protostellar binary IRAS~16293--2422.} {We use the data from an unbiased molecular line survey of the protostellar binary IRAS 16293$-$2422 between 329 and 363~GHz from the Atacama Large Millimeter/submillimeter Array (ALMA). The data probe scales of 60~AU (diameter) where most of the organic molecules are expected to have sublimated off dust grains and be present in the gas-phase. The deuterated and $^{13}$C-isotopic species of ketene, acetaldehyde and formic acid, as well as deuterated ethanol, are detected unambiguously for the first time in the interstellar medium. These species are analysed together with the $^{13}$C isotopic species of ethanol, dimethyl ether and methyl formate along with mono-deuterated methanol, dimethyl ether and methyl formate. } {The complex organic molecules can be divided into two groups with one group, the simpler species, showing a D/H ratio of $\approx$~2\% and the other, the more complex species, D/H ratios of 4--8\%. This division may reflect the formation time of each species in the ices before or during warm-up/infall of material through the protostellar envelope. No significant differences are seen in the deuteration of different functional groups for individual species, possibly a result of the short time-scale for infall through the innermost warm regions where exchange reactions between different species may be taking place. The species show differences in excitation temperatures between 125~K and 300~K. This likely reflects the binding energies/sublimation temperatures of the individual species, in good agreement to what has previously been found for high-mass sources. For dimethyl ether the $^{12}$C/$^{13}$C ratio is found to be lower by up to a factor of 2 compared to typical ISM values similar to what has previously been inferred for glycolaldehyde. Tentative identifications suggest that the same may apply for $^{13}$C isotopologues of methyl formate and ethanol. If confirmed, this may be a clue to their formation at the late prestellar/early protostellar phases with an enhancement of the available $^{13}$C relative to $^{12}$C related to small differences in binding energies for CO isotopologues or the impact of FUV irradiation by the central protostar.} {The results point to the importance of ice surface chemistry for the formation of these complex organic molecules at different stages in the evolution of embedded protostars and demonstrate the use of accurate isotope measurements for understanding the history of individual species.} | The earliest stages of protostars are characterised by a rich chemistry, which in some cases leads to abundant complex organic, and even prebiotic, molecules present in the gas-phase close to the young stars of both low- and high-mass. The general picture is that the formation of icy grain mantles and chemistry in and on grains facilitates the build-up of larger species through surface reactions starting from CO \citep[e.g.,][]{tielens82,hasegawa93,charnley97,watanabe02,fuchs09,oberg09,garrod13,fedoseev15,chuang17}. However, there are still a number of important open questions concerning the chemistry leading to the formation of these species, for example, what specific reactions are dominant, what is the importance of, e.g., UV irradiation, how important is the ice composition and the rate at which the ices are warming up. The isotopic composition of the gas is a particular characteristic of the early stages of young stars, which may hold significant clues to some of these networks. It has for some time been recognised that the many species in the cold gas around embedded protostars are significantly enhanced in deuterium relative to hydrogen \citep[e.g.,][]{vandishoeck95,ceccarelli01d2co} compared to the cosmic D/H ratio of $1.5-2.0\times 10^{-5}$ \citep[e.g,.,][]{linsky03,prodanovic10}. In particular, the deuterated isotopologues of species such as formaldehyde (H$_2$CO) and methanol (CH$_3$OH) are found to be significantly enhanced up to levels of $\sim$~10\% or above compared to the regular isotopologues in single-dish observations \citep[e.g.,][]{parise06deuterium} and some of these sources even show detections of multiply deuterated species (e.g., CD$_2$HOH and CD$_3$OH; \citealt{parise02} and \citealt{parise04}) indicating very high degrees of deuteration \citep[see also overview in][]{jorgensen16}. These enhancements are thought to be a consequence of the exothermic reaction \begin{equation} {\rm H}_3^+ + {\rm HD} \quad\leftrightarrows\quad {\rm H}_2{\rm D}^+ + {\rm H}_2+ \Delta E \end{equation} (where $\Delta E = 232$~K), which followed by dissociative recombination of H$_2$D$^+$ with electrons leads to an enhanced atomic D/H ratio that can be transferred to grains when the neutral species freeze-out \citep{tielens83}. Other isotopic systems such as $^{12}$C and $^{13}$C-containing isotopologues may also show variations compared to the local ISM due to the fractionation processes in the cold phases although the variations of those are expected to be at much smaller levels than those of deuterium vs. hydrogen. Isotope selective photodissociation or ion-molecule reactions in the gas through the reaction: \begin{equation} {\rm ^{13}C}^+ + {\rm CO}\quad \leftrightarrows \quad {\rm C}^+ + {\rm^{13}CO}+\Delta E \end{equation} (where $\Delta E = 35$~K; \citealt{langer84,furuya11}) may enhance the $^{13}$CO abundance in the gas-phase relative to $^{12}$CO. Isotope selective photodissociation may decrease the amount of $^{13}$CO relative to $^{12}$CO, but at the same time enhancing the amount of $^{13}$C. Once on the grains, the enhanced $^{13}$CO or $^{13}$C can then be incorporated into the complex organic molecules. Other mechanisms, such as segregation between $^{12}$CO and $^{13}$CO in the ices due to slight differences in their binding energies \citep[e.g.,][]{smith15} may also affect the amount of $^{12}$C relative to $^{13}$C that is available for incorporation into larger molecules. Since all of these fractionation processes are very sensitive to the gas physics (density and temperature) it has therefore been suggested that measurements of the relative abundances of different isotopologues of individual molecules may hold strong clues to the formation pathways constrained by chemical models or alternatively to the physical conditions and early evolution of young stars through their pre- or protostellar stages \citep[e.g.][]{charnley04,cazaux11,taquet13,taquet14}. With the Atacama Large Millimeter/submillimeter Array (ALMA) it is becoming possible to push forward in providing accurate measurements of the isotopic composition of complex organics. As an example of such a study, \cite{belloche16} presented a systematic survey of the deuteration of molecules toward the Sagittarius B2 high-mass star forming region in the Central Molecular Zone close to the Galactic Center. Utilising ALMA's high sensitivity, \citeauthor{belloche16} presented new detections as well as strict upper limits for the deuterated variants of a number of complex organics. Typically the abundances of the deuterated species were found to be lower than what is seen in nearby molecular clouds and also predicted by astrochemical models. \cite{belloche16} speculated that this may reflect a lower overall deuterium abundance or the higher temperatures in the clouds toward the Galactic Center. One of the prime targets for similar studies of solar-type stars is the Class~0 protostar IRAS~16293--2422. This source has long been recognised as having large abundances of deuterated species compared to high-mass protostellar sources \citep{vandishoeck95} and was also the first solar-type protostar for which abundant complex organic molecules were found \citep{cazaux03}. Combining these two properties, it has been a natural target for deuteration studies. This is the source toward which the above mentioned multi-deuterated species, as well as some deuterated complex organics, e.g., methyl formate (CH$_3$OCHO; \citealt{demyk10}) and dimethyl ether (CH$_3$OCH$_3$; \citealt{richard13}), were first detected. Recent observations of the THz ground state transitions of H$_2$D$^+$ \citep{bruenken14} and D$_2$H$^+$ \citep{harju17} toward the cold envelope around IRAS~16293--2422 suggest very similar amount of both cations in this source. Due to sensitivity limitations, however, the single-dish observations typically targeted the brightest, low excitation, transitions that in some cases may be optically thick both for the main isotopologues as well as for the deuterated variants. Furthermore, single-dish observations naturally encompass the entire protostellar envelope and are thus intrinsically biased toward emission on larger scales where most of the material is located and molecules may still be largely frozen-out on dust grains due to the lower temperatures. Using ALMA we have recently completed a large unbiased survey, the \emph{Protostellar Interferometric Line Survey (PILS)}, of a key frequency window around 345~GHz \citep{jorgensen16}. These data are up to two orders of magnitude more sensitive than previous single-dish studies and thus provide excellent opportunities for systematic studies of complex organics and their isotopologues. In particular, the interferometric observations make it possible to zoom in on the inner 30--60~AU around the central protostars, i.e., Solar System scales, where the temperatures are high and the ices have sublimated and the complex organic molecules therefore are all present in the gas-phase. The initial results presented the first detections of the deuterated isotopologues of isocyanic acid and formamide \citep{coutens16} as well as the deuterated and $^{13}$C-isotopologues of glycolaldehyde \citep{jorgensen16}. These species all show significant deuterium enhancements -- although not on the levels hinted by the previous single-dish observations -- but also indications of differences that may exactly reflect their chemistries. Specifically, glycolaldehyde is more enhanced in deuterium than the other species by a factor of a few. The $^{13}$C-isotopologues of glycolaldehyde further indicate an enhancement of the $^{13}$C isotopologues above that of the local ISM, again possibly reflecting formation of this species in the coldest part of the envelope. In the present work, we present an extensive investigation into the deuteration and the $^{13}$C content of oxygen-bearing complex organic molecules. These include methanol, its next heavier homologue ethanol along with its isomer dimethyl ether, the dehydrogenated relatives of ethanol, acetaldehyde, and ketene, as well as methyl formate, an isomer of glycolaldehyde, and formic acid. The paper is laid out as follows: Sect.~\ref{analysis} summarises the main points about the ALMA data and spectroscopy while Sect.~\ref{results} describes the data analysis, including line identification and column density/abundance determinations of the complex organics and their isotopologues. Sect.~\ref{discussion} discusses the implications in relation to the formation of the species. | We have presented a systematic survey of the isotopic content of oxygen-bearing (and some nitrogen-bearing from \citealt{coutens16}) complex organic molecules toward the low-mass protostar IRAS16293B. Using data from the ALMA \emph{Protostellar Interferometric Line Survey (PILS)} program we provide an inventory of those species and their relative abundances in the warm gas close to the central protostar where ices are completely sublimated. The main conclusions are: \begin{itemize} \item For the first time we identify the deuterated and $^{13}$C-isotopic species of ketene, acetaldehyde and formic acid as well as deuterated ethanol in the interstellar medium and present observations of $^{13}$C isotopic species of dimethyl ether, methyl formate and ethanol (the latter two only tentatively detected) along with mono-deuterated methanol, dimethyl ether and methyl formate. Systematic derivations of excitation temperatures and column densities result in small uncertainties in their relative abundances. \item Some differences are found in excitation temperatures for different species, with the lines for one group (formaldehyde, dimethyl ether, acetaldehyde and formic acid, together with formaldehyde, \citealt{persson18}), best fit with an excitation temperature of approximately 125~K with the remaining species, together with formamide and glycolaldehyde \citep{coutens16,jorgensen16}, better fit with a temperature of 300~K. This possibly reflects different binding energies of the individual species to the grains with the former group having binding energies in the 2000--4000~K range and the latter in the 5000--7000~K range. The division is very similar to what has previously been found for high-mass sources \citep{bisschop07}. \item The D/H ratios of the complex organic molecules can be divided into two groups with some of the simpler species (methanol, ketene, formic acid) as well as formaldehyde, formamide and isocyanic acid) showing D/H ratio of $\approx$~2\% and the more complex species dimethyl ether, ethanol, methyl formate, glycolaldehyde and acetaldehyde showing higher ratios ranging from about 4 to 8\%. Conservative estimates of the errors on these ratios are 0.6 percentage points for the lower value and 1.5 percentage points for the higher values. This distinction may reflect the formation time of each species in the ices before or during warm-up/infall of material through the protostellar envelope. \item No significant differences are seen in the deuteration of different functional groups for individual species, possibly a result of the short time-scale for infall through the innermost regions where exchange reactions between different species may be taking place. \item The $^{12}$C/$^{13}$C ratio of dimethyl ether is found to be lower than that of the local ISM, similar to the case of glycolaldehyde \citep{jorgensen16}. Marginal detections of the $^{13}$C isotopologues of methyl formate and ethanol are also consistent with a lower ratio. Low $^{12}$C/$^{13}$C ratios may reflect the formation histories of individual species; that they form in the ices at a late point where more $^{13}$C is available due to a slightly lower binding energy of $^{12}$CO compared to $^{13}$CO ice -- or fractionation triggered by FUV irradiation from the central protostar. \end{itemize} This study is an important illustration of how the isotopic composition can be used to determine the formation pathways for different molecular species. Efforts on analysing the nitrogen-containing organics toward IRAS16293B \citep{ligterink17,ligterink18,coutens18,calcutt18}, as well as establishing comparable inventories of both oxygen- and nitrogen-bearing species toward IRAS16293A, are ongoing. Together these studies will serve as an important constraint on astrochemical models attempting to account for the fractionation processes in protostellar environment. It would be worthwhile revisiting the abundances of species on larger scales for IRAS16293 and other sources, given the issues with opacity effects for many of the main isotopologues. As noted above, additional efforts are needed to investigate any deviations of the $^{12}$C:$^{13}$C ratios for the organics that are only marginally detected or those suffering from optical depth issues. Finally, extending this survey to other embedded protostars, possibly in other regions, could shed light on the importance of environment versus evolutionary histories of the sources for the formation of complex molecules. | 18 | 8 | 1808.08753 |
1808 | 1808.04057_arXiv.txt | The remnant radio galaxies in galaxy clusters are important sources of seed relativistic electron population in the intra-cluster medium (ICM). Their occurrence and spectral properties are poorly studied. In this work we present a broadband study of the radio relic in the galaxy cluster Abell 4038 using the Upgraded Giant Metrewave Radio Telescope (uGMRT). We present the uGMRT images in the bands 300 - 500 MHz and 1050 - 1450 MHz having rms noise $70\,\mu$Jy beam$^{-1}$ and $30\,\mu$Jy beam$^{-1}$, respectively, that are the deepest images of this field so far. A spectral analysis of the relic over 300 - 1450 MHz using images in sub-bands scaled to have constant fractional bandwidths to achieve a closely matched uv-coverage was carried out. The 100 kpc extent of the relic is divided into Loop, Arc, Bridge and North-end. The Loop has a steep spectral index of $\alpha=2.3\pm0.2$ ($S_{\nu}\propto\nu^{-\alpha}$). The North-end has ultra-steep spectra in the range $2.4 - 3.7$. The Arc is found to skirt a curved region seen in the \emph{Chandra} X-ray surface brightness image and the highest spectral curvature in it reaches $1.6\pm0.3$. We interpret the morphology and spectral properties of the relic in the scenario of an adiabatically compressed cocoon from the past activity of the Brightest Cluster Galaxy in the cluster. A comparison of the properties of the A4038 relic with a sample of 10 such relics is discussed. | \label{intro} Clusters of galaxies contain among the largest pools of baryons in the Universe called the intra-cluster medium (ICM). The ICM contains thermal gas at temperatures $10^7 - 10^8$ K and is also known to be magnetized at $\mu$G levels. About a third of the most massive clusters, radio synchrotron sources of extents of Mpc, co-spatial with the X-ray emission from the thermal gas or at the periphery are found \citep[e. g. ][a review]{bru14}. These provide evidence for the presence of electrons with relativistic energies in the ICM. Since diffusion time of such electrons from a single source across the cluster is far longer than their radiative lifetimes, processes of re-acceleration are invoked \citep{jaf77}. The cluster-wide diffuse synchrotron sources called the radio halos are believed to originate from in-situ re-acceleration of electrons via turbulence injected during cluster mergers \citep[][]{1987A&A...182...21S,pet01,bru04,bru11,bru16}. The arc-like radio sources at cluster peripheries are proposed to originate in the re-acceleration of mildly relativistic electrons at cluster merger shocks \citep[e. g.][]{ens98}. The turbulent re-acceleration and diffusive shock acceleration at the weak shocks in the ICM are both inefficient processes requiring, a seed population of mildly relativistic electrons \citep[e. g.][]{Kang2007,Bruggen2012,pin13,pin17}. The sources of seed relativistic electrons could be the secondary electrons produced in the hadronic collisions \citep{bru11} or those injected by the cluster galaxies \citep{bru14}. The Active Galactic Nuclei (AGN) and radio galaxies inject relativistic plasma via jets during the active phase of the central supermassive blackhole. However the jets and lobes of radio galaxies will fade on the timescales of $10^7-10^8$ years unless there are mechanisms at work that re-energize them \citep[e. g.][]{1987A&A...182...21S}. Adiabatic compression of remnant radio cocoons due to shocks and disturbances in the ICM has been proposed to revive the radio emission \citep{ens01}. Simulations of this process predict that the compressed cooon will be shred into filaments allowing dispersal of the plasma \citep{2002MNRAS.331.1011E}. A mechanism of gentle re-acceleration also has been proposed recently \citep{2017NatAs...1E...5V}. Such remnants that show revival from their fading phase have also been referred to as radio ``phoenixes'' \citep{kemp04} in the literature. In this work we focus on remnant relics that could be important sources of seed relativistic electrons in the ICM. The detection of such remnants in clusters is challenging due to their low brightness and rare occurrence due to short radiative lifetimes. A sample of 13 such relics has been presented in \citet{fer12} where these are classified as ``roundish'' relics to distinguish them from the arc-like and elongated relics associated with merger shocks. A further larger sample of ``phoenix'' candidates has been presented in \citet{nuza17} and recently a faint sub-sample of these relics has been presented by \citet{Wilber2018}. Integrated spectra of such relics over broad ranges of frequencies have been measured for some of the relics \citep{sle01,wee11b,cohen11,kaldwa12}. Spectral index studies of a few relics with good resolution across the relic have also been carried out by combining multi-frequency radio data from one or more radio telescopes and have revealed the complex morphologies and spectra and allowed to estimate the life cycle of the relativistic plasma in them \citep[e. g.][]{cohen11,kaldwa12,2015A&A...583A..89S,2017A&A...600A..65S}. Here we present a broadband study of a remnant radio relic towards the galaxy cluster Abell 4038 (A4038, hereafter) using the Upgraded Giant Metrewave Radio Telescope (uGMRT). The uGMRT observations resulting in the deepest images towards this source are presented and used to measure spectral curvature across the extent of the relic. The radio morphology, spectra and the X-ray surface brightness map of the cluster together point to a scenario in which the relic originated from the compression of a radio cocoon left by the central galaxy. A sample of such relics is presented and the A4038 relic in comparison with others is discussed and the importance of spatially resolved wideband observations of these sources is emphasized. The paper is organized as follows: Sec.~\ref{a4038} describes the cluster Abell 4038 and the relic in it. The observations and data analysis are described in Sec.~\ref{obs}. The uGMRT images are presented in Sec.~\ref{ugmrtim}. The spectral curvature analysis is described in Sec.~\ref{curv}. The results are discussed in Sec.~\ref{discussion}. The summary and conclusions are presented in Sec.~\ref{conclusions}. We have used $\Lambda$CDM cosmology with $H_0$=$70$ km s$^{-1}$ Mpc$^{-1}$, $\Omega_\Lambda$=$0.73$ and $\Omega_m$ = $0.27$ in this work. | We have presented a spectral curvature study of the relic in Abell 4038 using the uGMRT in the 300 - 500 MHz and 1050 - 1450 MHz bands. The connection between the relic properties and the X-ray emission from the cluster was discussed. A comparison between the A4038 relic and other small scale cluster relics was also carried out. The main findings are summarized below: \begin{enumerate} \item We have produced deep images of the radio relic in A4038 at 0.402 and 1.26 GHz using the uGMRT with bandwidths of 200 and 400 MHz in the two bands, respectively. The rms noise at the centre of the field at 402 MHz was $0.07$ mJy beam$^{-1}$ and at 1.26 GHz was $0.03$ mJy beam$^{-1}$. The largest extent of the relic in the north-south direction is 102 kpc and in the east-west direction is 58 kpc. \item The spectra across the relic were estimated in 15 regions, each of size $15''\times15''$. The spectrum of each region was fit with separate power-laws in the ranges 0.32 - 0.45 and 0.45 - 1.43 GHz. The difference in the spectral indices in the low ($\alpha_{\rm low}$) and high frequency ($\alpha_{\rm high}$) ranges was used as a measure of the spectral curvature ($\Delta \alpha$). \item The highest curvature, $\Delta \alpha$, of $1.6\pm0.3$ was found in the region corresponding to the Arc. The curvature is 0.8 in the regions 10 and 11 that make the eastern and southern parts of the Loop, respectively. To the west of the Loop it becomes consistent with zero in region 12 and then rises to 0.5 further west. \item A curved arc-like region in the X-rays is skirted by the Arc in the relic. We propose that the highly curved spectra in the Arc result from compression-revived radio emission. This is consistent with our earlier work presented in \citet{kaldwa12}, where the integrated spectrum was best fit in the adiabatically compressed phase in the model by \citet{ens01}. \item The calculation of magnetic field under the minimum energy condition is dependent on the spectral index and the current formalisms assume a single power-law. For A4038 relic if we assume a spectral index of 1.5, the magnetic field is 7.8 $\mu$G if the minimum Lorentz factor, $\gamma_{\rm min}=100$. \item We also present a sample of 10 remnant relics in galaxy clusters that have detectable X-ray emission from their ICM. These relics are typically found within the 0.5 Mpc from the cluster centre-the only exception being A1664. \item We have plotted low versus high frequency spectral indices for the remnant relic sample, together with the values for the A4038 relic. Except the A1664 relic which has a flat and straight spectrum, the sample remnant relics have curvatures in the range 0.5 - 1.6. These are consistent either with an old synchrotron plasma or a plasma re-energised by a mechanism such as adiabatic compression by shock or gentle re-acceleration. \item We conclude that for A4038 the spectral shape changes across the regions within the relic and thus spatially resolved spectral studies as presented for the case of A4038 relic are essential to study the origins of such relics. \item The deep radio images from the uGMRT presented here also resulted in the detection of faint lobes around a radio core identified with the galaxy PGC072471. \end{enumerate} | 18 | 8 | 1808.04057 |
1808 | 1808.09878_arXiv.txt | PLATO\cite{b1} has been selected and adopted by ESA as the third medium-class Mission (M3) of the Cosmic Vision Program, to be launched in 2026 with a Soyuz-Fregat rocket from the French Guiana. Its Payload (P/L) is based on a suite of 26 telescopes and cameras in order to discover and characterise, thanks to ultra-high accurate photometry and the transits method, new exoplanets down to the range of Earth analogues. Each camera is composed of 4 CCDs working in full-frame or frame-transfer mode. 24 cameras out of 26 host 4510 by 4510 pixels CCDs, operated in full-frame mode with a pixel depth of 16 bits and a cadence of 25 s. Given the huge data volume to be managed, the PLATO P/L relies on an efficient Data Processing System (DPS) whose Units perform images windowing, cropping and compression. Each camera and DPS Unit is connected to a fast SpaceWire (SpW) network running at 100 MHz and interfaced to the satellite On-Board Computer (OBC) by means of an Instrument Control Unit (ICU), performing data collection and compression. | The PLATO Payload, based on two groups of 24 and 2 telescopes, covering ~2250 square degrees per pointing, is conceived for the study of planetary systems formation and evolution and to answer fundamental questions concerning the existence of other planetary systems like our own, including the presence of new worlds in the habitable zone of Sun-like stars. The 24 telescopes, observing faint stars, and the 2 telescopes, observing bright targets, named respectively \emph{normal} and \emph{fast} telescopes, are operated by their front-end electronics (FEEs, each controlling 4 CCDs as Focal Plane Array -FPA- detectors) and will provide the capability to attain a large photometric visible magnitude range, from ~4 to ~16. Focusing on a subset of brighter targets (m\textsubscript{v} 4-11), the PLATO Payload will detect and characterise planets down to the Earth size by means of photometric transits and thanks to the determination of their masses through ground-based radial velocity follow-up campaigns. Given the brightness of this subset of samples, PLATO will extensively perform asteroseismology on these targets to retrieve highly accurate stellar parameters such as masses, radii and ages allowing for a precise characterisation of planetary bulk parameters. The main scientific requirement to detect and characterise a large number of terrestrial planets around bright stars plays a key role in defining the PLATO observing strategy and its own Payload. The current baseline observing plan for the 4-years nominal science operations consists of long-duration observations of two sky fields lasting two years each. An alternative scenario is for operations split into a long-duration pointing lasting three years and a one-year step-and-stare phase with several pointings. Long pointings will guarantee the detection of planets down to the habitable zone of solar-like stars with a first basic assessment of the main characteristics of their atmospheres, opening the way to future space missions (e.g. ARIEL) designed to perform spectroscopy on these targets. \begin{figure}[htbp] \centerline{\includegraphics[height=5cm]{PLATO_satellite.eps}} \caption{The PLATO satellite with its unique Payload composed by 26 small telescopes in an artistic vision.} \label{PLATO_DPS} \end{figure} The Payload Data Processing System (DPS) is composed of 2 sets of 6 Normal-DPUs (Data Processing Units), each one collecting data from two N-telescopes and 2 Fast-DPUs (one per each Fast camera/telescope) designed to pre-process the downloaded images extracting the photometric signal from the targets and selecting a reduced sample of imagettes (small images of $n$x$n$ pixels) around a subset of stellar samples. The DPUs will also be in charge of the detection and removal of outliers from the photometric signals. An Instrument Control Unit (ICU) will collect, thanks to the P/L SpW network, all the scientific data and payload housekeeping (HK) and will compress them before delivery to the Service Module (SVM), as it will act as the main interface between the Payload and the spacecraft (S/C). The ICU will also be in charge of the DPS SpaceWire network management. It is conceived as a full cold-redundant unit implementing a minimal cross-strapping to improve the DPS reliability. | The overall PLATO P/L design relies, as described in this paper, on an efficient cold redundant SpaceWire network. At the present time, thanks to several iterations between ESA, the PLATO Mission Consortium (PMC) and the Spacecraft Prime Contractor (OHB, Bremen), it is under study the possibility to switch-on, when in-flight, both the nominal and redundant sides at the same time in order to manage the remote likelihood to have a double failure on the SpW network and its components. This will be one of the major outcomes of the co-engineering phase between PMC and ESA/Prime, that will end by July, 2018 and will affect the ICU design, under the INAF responsibility. | 18 | 8 | 1808.09878 |
1808 | 1808.05098_arXiv.txt | The Earth is well-known to be, in the current astronomical configuration, in a regime where two asymptotic states can be realised. The warm state we live in is in competition with the ice-covered snowball state. The bistability exists as a result of the positive ice-albedo feedback. In a previous investigation performed on a intermediate complexity climate model we have identified the {\color{black}unstable climate} states (\textit{Melancholia states}) separating the co-existing climates, and studied their dynamical and geometrical properties. The Melancholia states are ice-covered up to the mid-latitudes {\color{black}and attract} trajectories initialised on the basins boundary. In this paper, we study {\color{black}how stochastically perturbing the parameter controlling the intensity of the incoming solar radiation impacts} the stability of the climate. We detect transitions between the warm and the snowball state and analyse in detail the properties of the noise-induced escapes from the corresponding basins of attraction. We {\color{black}determine} the most probable paths for the transitions and find evidence that the Melancholia states act as gateways, similarly to saddle points in an energy landscape. | 18 | 8 | 1808.05098 |
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1808 | 1808.02447_arXiv.txt | Despite progress in understanding radio relics, there are still open questions regarding the underlying particle acceleration mechanisms. In this paper we present deep 1--4 GHz VLA observations of CIZA\,J2242.8+5301 ($z=0.1921$), a double radio relic cluster characterized by small projection on the plane of the sky. Our VLA observations reveal, for the first time, the complex morphology of the diffuse sources and the filamentary structure of the northern relic. We discover new faint diffuse radio emission extending north of the main northern relic. Our Mach number estimates for the northern and southern relics, based on the radio spectral index map obtained using the VLA observations and existing LOFAR and GMRT data, are consistent with previous radio and X-ray studies ($\mathcal{M}_{\rm RN}=2.58\pm0.17$ and $\mathcal{M}_{\rm RS}=2.10\pm0.08$). However, color-color diagrams and modelings suggest a flatter injection spectral index than the one obtained from the spectral index map, indicating that projection effects might be not entirely negligible. The southern relic consists of five ``arms''. Embedded in it, we find a tailed radio galaxy which seems to be connected to the relic. A spectral index flattening, where the radio tail connects to the relic, is also measured. We propose that the southern relic may trace AGN fossil electrons that are re-accelerated at a shock, with an estimated strength of $\mathcal{M}=2.4$. High-resolution mapping of other tailed radio galaxies also supports a scenario where AGN fossil electrons are revived by the merger event and could be related to the formation of some diffuse cluster radio emission. | \label{sec:intro} Cluster mergers are the most energetic phenomena in the Universe since the Big Bang, involving kinetic energies of $\sim 10^{63}-10^{65}$ erg, released over a time scale of 1--2 Gyr \citep[e.g.][]{sarazin02} depending on the cluster's mass and on the relative velocity of the merging dark matter halos. In a hierarchical cold dark matter (CDM) scenario, these phenomena are the natural way to form rich cluster of galaxies. Cluster mergers produce shock waves and turbulence in the intracluster medium (ICM). It has been proposed that ICM shocks and turbulence can (re-)accelerate cosmic rays (CRs) which then produce diffuse radio synchrotron emitting sources in the presence of $\mu$Gauss magnetic fields. These diffuse sources are known as {\it radio halos} and {\it radio relics} \citep[see][for a theoretical and an observational review]{brunetti+jones14,feretti+12}. Radio halos are centrally located unpolarized sources with a smooth morphology, roughly following the X-ray emission from the ICM. The currently favored formation scenario for the formation of radio halos involves the re-acceleration of pre-existing CR electrons \citep[energies of $\sim$ MeV, e.g.][]{brunetti+01, petrosian01} via magneto-hydrodynamical turbulent motions, induced by a merging event. Moreover, radio halos are characterized by a steep-spectrum \citep[$\alpha < -1$, with $S_\nu\propto\nu^\alpha$, e.g.][]{brunetti+08}. A correlation between the halo radio power and the cluster X-ray luminosity exists \citep[e.g.][]{cassano+13}, showing that the most X-ray luminous clusters host the most powerful radio halos. Moreover, \cite{cassano+10} showed a clear correlation between the cluster's dynamical state, measured from the X-ray surface brightness distribution, and the presence of giant radio halos, providing strong support that mergers play an important role in the formation of these sources. As with the radio halos, radio relics are characterized by a steep spectral index ($\rm \alpha\approx-0.8~to~-1.5$). They are generally elongated sources located in the outskirts of clusters, and they are suggested to trace outwards traveling shock fronts, where the ICM and magnetic fields are compressed. \citep[e.g.][]{ensslin+98,finoguenov+10,vanweeren+10}. As a consequence, the magnetic field is amplified and aligned, producing polarized radio emission. One formation scenario for relics is the {\it diffusive shock acceleration} (DSA) mechanism, similar to what occurs in supernova remnants \citep[e.g.][]{blandford+ostriker78,drury83,ensslin+98}. This model involves particles (electrons) that are accelerated from the ICM's thermal pool into CRs at shocks, while the electrons in the downstream region suffer synchrotron and inverse Compton (IC) energy losses. A prediction of the DSA model is a relation between the Mach number of the shock and the radio spectral index at the shock location \citep[e.g.][]{giacintucci+08}. However, recent observations suggest fundamental problems with the standard DSA model: (1) in a few cases, there are cluster merger shocks without corresponding radio relics \citep[e.g. the main shock in the Bullet Cluster,][]{shimwell+14} (2) sometimes, the spectral index derived Mach numbers are significantly higher than those obtained from the X-ray observations \citep[e.g.][]{vanweeren+16}, and (3) some relics require an unrealistic shock acceleration efficiency (for DSA) to explain their observed radio power \citep[e.g.][]{vazza+14,botteon+16}. Another formation scenario is {\it shock re-acceleration} of relativistic fossil electrons \citep[e.g.][]{markevitch+05, macario+11, kang+ryu11, kang+12, bonafede+14, shimwell+15, kang+17, vanweeren+17a}. This mechanism addresses the DSA efficiency problem, and implies a direct connection between radio relics and radio galaxies, which are supposed to be the primary sources of fossil plasma. Radio ``phoenices'' are another class of relics that are characterized by their steep curved spectra and often complex toroidal morphologies. They are supposed to trace AGN fossil plasma lobes which have been adiabatically compressed by merger shock waves \citep{ensslin+gopal-krishna01,ensslin+bruggen02,degasperin+15}. Relics have a range of sizes and morphologies. So-called {\it double relic} systems, with two relics located on diametrically opposite sides of the cluster center \citep[e.g.][]{rottgering+97, bagchi+06, venturi+07, bonafede+09, vanweeren+09, bonafede+12, degasperin+14} are an important subclass. Numerical simulations \citep[e.g.][]{vanweeren+11} suggest that these systems are the result of binary mergers between two-comparable mass clusters (mass ratio of 1:1 or 1:3), with the line connecting the two relics representing the projected merger axis of the system. In this paper we present a radio continuum and spectral analysis of the merging cluster CIZA\,J2242.8+5301, which hosts two opposite radio relics and a faint radio halo, by means of new deep 1--4~GHz Karl G. Jansky Very Large Array (VLA) observations. The paper is organized as follows: in Sect. \ref{sec:sausage} we provide an overview of CIZA\,J2242.8+5301, presenting the results of previous studies performed in the optical, X-ray and radio bands; in Sect. \ref{sec:obs} we describe the radio observations and the data analysis; the radio images, the spectral index maps are presented Sect. \ref{sec:results}. We end with a discussion and conclusions in Sects. \ref{sec:discuss} and \ref{sec:concl}, respectively. In this paper, we adopt a flat $\Lambda$CDM cosmology with H$_0=70$ km s$^{-1}$ Mpc${-1}$, $\Omega_{\rm m}=0.3$ and $\Omega_\Lambda=0.7$, which implies a conversion factor of 3.22 kpc/$^{\prime\prime}$ and a luminosity distance of $\approx944$ Mpc, at the cluster's redshift \citep[$z =0.192$][]{kocevski+07}. | \label{sec:concl} We have presented deep 1--4 GHz VLA observations of the galaxy cluster CIZA\,J2242.8+5301 ($z=0.1921$). This cluster is one of the best examples of a merging system, hosting two Mpc-size relics and several other diffuse radio sources. Our images reached a resolution of $2.1^{\prime\prime}\times1.8^{\prime\prime}$ and $0.8^{\prime\prime}\times0.6^{\prime\prime}$ and a noise level of $3.8~\mu$Jy~beam$^{-1}$ and $2.7~\mu$Jy~beam$^{-1}$ at 1.5 and 3.0~GHz respectively. These observations were combined with existing GMRT 610 MHz and LOFAR 150 MHz data \citep{vanweeren+10,stroe+13,hoang+17} to carry out a radio continuum and spectral study of the cluster. The main results of our study are summarized below: \begin{itemize} \item[$\bullet$] The high-resolution images reveal that the northern relic (RN) is not a continuous structure, but is broken up into several filaments with a length of $\sim 200-600$ kpc each. Moreover, we detect and characterize additional radio emission north of RN, labeled as R5. \medskip \item[$\bullet$] In agreement with previous studies, we find a trend of North-South spectral steepening for the northern relic, indicative of IC and synchrotron energy losses in the downstream region of a shock. We measure an injection spectral index at the shock front from the radio map of $-0.86\pm0.05$, corresponding to a Mach number of $\mathcal{M}=2.58\pm0.17$. Although this value is consistent with the Mach number estimated from X-ray studies, a color-color diagram shows a small amount of spectral curvature at the relic's northern edge indicating mixing of emission with different spectral ages, possibly due to projection effects. This suggests that the true injection spectral index is slightly flatter ($\alpha\sim-0.7$). \medskip \item[$\bullet$] The low-resolution VLA images reveal that the southern relic (RS) has a complex morphology, characterized by the presence of five ``arms''. The origin of this complex morphology is unclear, but it might be partly explained by the presence of AGN fossil plasma that has been revived by the passage of a shock wave. Interestingly, the color-color diagram revealed no curvature in the downstream region of the southern relic, which can be explained either by a sum of power-law spectra with different $\alpha_{\rm inj}$ at different locations or smearing of broad range of magnetic fields and/or electron loss histories. \medskip \item[$\bullet$] We identify source~J, a tailed radio galaxy embedded in RS, as possible source of fossil radio plasma. Our spectral index maps suggest that old plasma from the lobe of source~J is re-accelerated. We attempted to model the flattening of the spectral index from the lobe to the relic RS. We find that a $\mathcal{M}=2.4$ is compatible with the observed spectral index change. \medskip \item[$\bullet$] The cluster contains several tailed radio galaxies, which show disturbed morphologies likely due to the interaction between the radio plasma, ICM, and the merger event. In addition, the cluster contains a number of complex diffuse radio sources that seem to be related to plasma from tailed radio sources (source~G, RS, I, R2). The range in morphologies of tailed-radio sources, strong signs of interactions with the ICM, and presence of nearby complex-shaped relics suggest that fossil plasma from cluster radio galaxies plays an important role in the formation of (at least some) diffuse radio sources. To produce realistic models of diffuse cluster radio emission, this fossil plasma will need to be included in simulations. \medskip \end{itemize} | 18 | 8 | 1808.02447 |
1808 | 1808.07781_arXiv.txt | With the arrival of a number of wide-field snapshot image-plane radio transient surveys, there will be a huge influx of images in the coming years making it impossible to manually analyse the datasets. Automated pipelines to process the information stored in the images are being developed, such as the LOFAR Transients Pipeline, outputting light curves and various transient parameters. These pipelines have a number of tuneable parameters that require training to meet the survey requirements. This paper utilises both observed and simulated datasets to demonstrate different machine learning strategies that can be used to train these parameters. We use a simple anomaly detection algorithm and a penalised logistic regression algorithm. The datasets used are from LOFAR observations and we process the data using the LOFAR Transients Pipeline; however the strategies developed are applicable to any light curve datasets at different frequencies and can be adapted to different automated pipelines. These machine learning strategies are publicly available as {\sc Python} tools that can be downloaded and adapted to different datasets (https://github.com/AntoniaR/TraP\_ML\_tools). | Transient and variable astronomy is entering an exciting era with vast numbers of images, covering large fields of view, requiring real-time automated processing. In radio astronomy in particular, new instruments including LOFAR \citep[the Low Frequency Array;][]{vanHaarlem:2013}, AARTFAAC \citep[Amsterdam-ASTRON Radio Transients Facility And Analysis Centre;][]{Prasad:2012}, MWA \citep[Murchison Wide-field Array;][]{Tingay:2013}, LWA1 \citep[Long Wavelength Array Station 1;][]{Ellingson:2013}, ASKAP \citep[Australian Square Kilometre Array Pathfinder;][]{Johnston:2007,Murphy:2013} and MeerKAT \citep{Booth:2009} are exploring unfamiliar frequency and transient duration regimes, opening up the possibility of discovering new types of variable sources. Significant effort is being applied to developing high quality tools that are capable of rapidly and reliably processing images to obtain light curves of all the sources in the fields imaged and to automatically analyse the light curves to find transient or variable sources\footnote{For clarity purposes, we define ``transient sources'' as those which are new sources detected after the first image in a given dataset and ``variable sources'' as those which have variability in their extracted light curves.}. The LOFAR Transients Pipeline \citep[{\sc TraP};][]{Swinbank:2014} is one such tool that is capable of processing images to find transient and variable sources from a wide range of radio telescopes. To date, TraP has been used by LOFAR, AARTFAAC, the Boolardy Engineering Test Array \citep[BETA, the ASKAP test array;][]{Hotan:2014,Hobbs:2016}, the Arcminute Microkelvin Imager \citep[AMI;][]{Zwart:2008,Staley:2013}, the MWA \citep[][]{Rowlinson:2016} and the Jansky Very Large Array Low-band Ionosphere and Transient Experiment \citep[VLITE;][]{Polisensky:2016}. Once light curves are obtained and analysed by automated pipelines, there are two key challenges: distinguishing transient and variable sources from stable sources and classification of these sources into known source types. A number of existing transient surveys, at a range of wavelengths, have made good progress in determining key diagnostic parameters for identifying transient sources. One of the most powerful diagnostics is the weighted reduced $\chi^2$, referred to as $\eta$ in this paper, of a fit to the light curve by a constant flux density model, which can easily be converted to a probability that the data are drawn from the fitted model \citep[currently used to search for transients at many wavelengths, including X-ray, optical, microwave and radio; e.g.][]{Bannister:2011, Bower:2011, Palanque:2011, Thyagarajan:2011, Hoffman:2012, Shin:2012, Chen:2013, Croft:2013, Mooley:2013, Williams:2013, Bell:2014, Franzen:2014, Bell:2015}. Many surveys utilise a probability threshold to separate the stable sources from the variable sources. For instance, transient sources have a probability of less than 1\% of being a stable source using the $\eta$ from the fit to a stable source, enabling thresholds to be determined by using a False Detection Rate (FDR, the number of false identifications that are expected to be made). Several other parameters are also used by transient surveys to characterise the variable sources after identification or with an arbitrary threshold to aid in discriminating between variable and stable sources. These parameters can be source specific, such as the pulsar modulation index \citep[e.g.][]{Esamdin:2004}, or target a specific type of variability behaviour, such as periodicity in the light curve. One of the most common additional parameters is the fractional variability \citep[also known as fractional modulation; e.g.][]{Bannister:2011, Bower:2011, Chen:2013, Croft:2013, Mooley:2013, Bell:2014, Franzen:2014, Bell:2015}, which compares the observed scatter in flux densities to the average flux density of the source. However, variable sources are rarely identified using the information available in these additional parameters; typically they are only used to characterise known sources. In some cases, multiple thresholds are used to discriminate between stable and variable sources but these thresholds are typically chosen arbitrarily based on knowledge of specific datasets. Therefore, there is additional information available to aid in identifying variable sources but currently not used to its full potential. Over recent years, machine learning strategies have been increasingly applied to astronomical datasets. Machine learning is a form of artificial intelligence in which computer algorithms learn from data, in either a supervised or unsupervised manner, producing models that can then be applied to new data. For instance, one of the common machine learning applications is for classification. The separation of variable sources (ideally into different types) and stable sources using properties of their light curves is a classification problem and, hence, there are a number of machine learning techniques that could be applied to the datasets. Machine learning algorithms exploit various features about each source that can then be utilised in the algorithm; in the context of variable sources these parameters could include $\eta$ and the fractional variability. The algorithm ``learns'' how to separate the different groups of sources using a training dataset and the algorithm can then be applied within automated pipelines to enable rapid classification of new data points. There are two key types of algorithms in machine learning that can be used: supervised learning and unsupervised learning. There are particular types of unsupervised learning algorithms that would be a good strategy to find unusual sources in the data, in which the algorithm automatically finds new variable sources within a discovery space and classifies sources with similar properties for later identification by astronomers. However, unsupervised techniques, including some powerful unsupervised deep learning algorithms, often require large amounts of data to train in scenarios such as that outlined in this paper. Here, the populations of transient and variable sources are very rare compared to the stable source population and hence a large dataset is needed to ensure a sufficiently large population of the target sources. The required dataset size required for training unsupervised strategies is very difficult to quantify as it depends upon the number of discrete classification groups (the number of which is unknown), having sufficient data in each expected classification group and the distinctness of each group given the training parameters available. When starting a new transient survey searching for rare sources, these sufficiently large training datasets are not available. Here, as outlined later, we create a labelled dataset with simulated sources making this problem well suited for supervised machine learning strategies. In future surveys, when these issues have been addressed, algorithms such as unsupervised deep learning are likely to prove useful. This is because there will be a large amount of parameters for each unique source and a powerful unsupervised algorithm may identify previously undiscovered transient or variable behaviour of the sources. Alternatively, supervised machine learning strategies can be trained to find and classify different types of sources. These strategies use training datasets that contain pre-classified sources, which can either be from manual identification of sources in real data or from simulated datasets. For instance, supervised machine learning strategies can be used to reliably discriminate between new sources and imaging artefacts \citep[e.g.][]{Bloom:2012, Brink:2013}. Additionally, a supervised machine learning algorithm, known as a Random Forest \citep{Breiman:2001}, has been trained to identify different types of optical sources using data from the Palomar Transients Factory \citep[PTF;][]{Bloom:2012}. A Random Forest technique has also been applied to 2XMM and 3XMM data \citep[the second and third releases of the X-ray serendipitous source catalogue from the {\it XMM-Newton} satellite;][]{Jansen:2001,Watson:2009,Rosen:2016} to enable optimal identification of specific types of sources \citep{Lo:2014a,Lo:2014b,Farrell:2014} and to the classification of fast radio bursts \citep[e.g][]{Wagstaff:2016}. The Random Forest algorithm is also being investigated as a classification tool within the VAST (Variable and Slow Transient) survey which will be conducted using ASKAP \citep{Murphy:2013}. However, supervised learning techniques are really good at identifying objects that fit into known classes, but not good at identifying objects that belong to new classes. Unfortunately, when exploring a completely new parameter space, this information is often unknown or speculative. In this work, we instead focus on the simpler classification of sources into the categories of transient, variable or stable, which can be completed without knowing the behaviour of specific transient sources. Alternatively, some studies are now looking at unsupervised machine learning techniques such as unsupervised deep learning algorithms \citep[e.g.][]{Connor:2018}. In this paper, we use the {\sc TraP} to analyse variability of sources within large datasets from LOFAR, both observational and simulated. The transient surveys conducted by LOFAR are searching a relatively unexplored parameter space in radio astronomy: wide fields of view (enabling detection of the rarest variable sources) at low frequencies and on a range of time scales. Hence, this new parameter space may contain previously unidentified categories of variable sources, so we do not have the labelled datasets required to train machine learning strategies such as Random Forest algorithms to find these new sources. Instead, we utilise various variability parameters, output by {\sc TraP}, combined with machine learning techniques, to determine thresholds and algorithms that can optimally separate variable sources from stable sources. We create a simple labelled dataset that can be used to train and test the algorithms used. In Section 2, we describe the LOFAR datasets used in developing these new techniques. Section 3 presents analyses of the variability parameters output from {\sc TraP} for the datasets. The implementation of machine learning strategies is explained in Section 4, including a simple anomaly detection algorithm, as proposed by \cite{Denning:1987}, that outputs thresholds suitable for use in {\sc TraP} and a penalised logistic regression algorithm \citep[e.g.][]{Darroch:1972} that uses multiple parameters to categorise datasets. Finally, in Section 5, we discuss future improvements to these strategies. | Reliable automated transient and variable source detection is essential with the expected influx of data from large transient surveys. Tools such as {\sc TraP} are essential for processing these datasets but require the astronomer to choose appropriate thresholds. This paper has presented methods for choosing these thresholds using machine learning strategies applied to observed and simulated datasets. Using an anomaly detection strategy, adapted to be a supervised machine learning strategy in 2 dimensions, we have developed a method for astronomers to choose thresholds to meet specified criteria on the precision and recall of the classification of variable sources. This enables flexibility in the choice of thresholds; for instance the precision can be increased when requiring extremely reliable triggers to carefully control choice of follow-up targets. Alternatively, the recall can be set to higher values when the user wants to maximise the number of candidate variable sources. With reasonable input values for precision and recall, we blindly detected two known scintillating pulsars in the LOFAR RSM dataset. However, although the anomaly detection strategy performs better than the arbitrary thresholds chosen in previous studies, we note that this strategy is slow to train and provides sub-optimal classifications in comparison to the other methods presented in this paper. The second strategy presented uses a logistic regression strategy that divides datasets into two populations in a multiple parameter space. We have used 4 parameters ($\eta_{\nu}$, $V_{\nu}$, maximum flux density and the ratio between the maximum flux density and average flux) but the algorithm can be easily adapted to include many more parameters. This method achieves a precision of 98\% and recall of 91\% with the dataset presented in this paper. Further analysis with larger datasets and a wider range of variable sources will be highly beneficial in further testing and extending this strategy. Although transient sources will often have high values for their variability parameters, this is not always the case as their light curves may have insufficient data points to confirm variability or, after they have appeared, the sources may have a reasonably stable flux. {\sc TraP} has been specifically designed to deal with these sources by checking if they would be detected in the previous best images. To enable fine tuning, a margin can be applied to the detection threshold and we have demonstrated strategies to determine the optimal margin for the dataset. These strategies have been developed using the {\sc TraP} and LOFAR observations. We note that the performance of the machine learning algorithms needs to be tested with significantly larger and varied datasets as they become available over the coming years. Moreover, these methods have been designed to be adaptable to other pipelines and wavelengths. In future work, we plan to extend these strategies and apply them to a wide range of transient and variable sources. | 18 | 8 | 1808.07781 |
1808 | 1808.07978_arXiv.txt | 4C +49.22 is a $\gamma$-ray flat spectrum radio quasar with a bright and knotty jet. We investigate the properties of the core and large-scale knots by using their spectral energy distributions (SEDs). Analyzing its \emph{Fermi}/LAT data in the past 8 years, a long-term steady $\gamma$-ray emission component is found besides bright outbursts. For the core region, the $\gamma$-ray emission together with the simultaneous emission in the low-energy bands at different epochs is explained with the single-zone leptonic model. The derived magnetization parameters and radiation efficiencies of the radio-core jet decrease as $\gamma$-ray flux decays, likely indicating that a large part of the magnetic energy is converted to the kinetic energy of particles in pc-scale. For the large-scale knots, their radio--optical--X-ray SEDs can be reproduced with the leptonic model by considering the inverse Compton scattering of cosmic microwave background photons. The sum of the predicted $\gamma$-ray fluxes of these knots is comparable to that observed with LAT at $\sim10^{24}$ Hz of the steady $\gamma$-ray component, indicating that the steady $\gamma$-ray emission may be partially contributed by these large-scale knots. This may conceal the flux variations of the low-level $\gamma$-ray emission from the radio-core. The derived bulk Lorentz factors of the knots decrease along the distance to the core, illustrating as deceleration of jet in large-scale. The powers of the core and knots are roughly in the same order, but the jet changes from highly magnetized at the core region into particle-dominated at the large-scale knots. | % \label{sect:intro} The substructures of large-scale jets in radio-loud active galactic nuclei (AGNs), i.e., knots, hotspots, and lobes, have been resolved at the radio, optical and X-ray bands (see Harris \& Krawczynski 2006 for a review). This presents an opportunity to reveal the jet properties from the radio-core to the large-scale knots. This is helpful for revealing jet formation and propagation, composition, particle acceleration and radiation mechanisms, etc. (e.g., Zargaryan et al. 2017). The very long baseline interferometry (VLBI) observations in multi-epoch measurements of sub-parsec scale jets suggest that AGN jets start out highly relativistic with a Lorentz factor of tens (Jorstad et al. 2005; Lister et al. 2016), and it was proposed that the jets are still mildly relativistic at the kpc-scale (e.g., Arshakian \& Longair 2004; Mullin \& Hardcastle 2009). Convincing evidence for jet deceleration and transverse motions in M87 is presented by measuring its kpc-scale proper motions with Hubble Space Telescope (HST) and the pc-scale proper motions with VLBI, and the apparent velocity that is still superluminal in its outer jet (Meyer et al. 2017). The broadband spectral energy distributions (SEDs) of both the radio-core and large-scale jet radiations for radio-loud AGNs shows that they are non-thermal emission origin and show a bimodal feature (e.g., Sikora et al. 1994; Ghisellini et al. 2009; Zhang et al. 2010, 2013, 2014, 2018). The high energy radiation beyond the X-ray band of the core region should be dominated by the synchrotron-self-compton scattering (SSC, Sikora et al. 1994; Ghisellini et al. 2009; Zhang et al. 2013) process and/or the inverse scattering the photons of the broad-line region (IC/BLR, Ghisellini et al. 2009; Zhang et al. 2014, 2015) or torus (IC/torus, Sikora et al. 2009; Kang et al. 2014). The high energy (the X-ray--$\gamma$-ray bands) radiation mechanisms of large-scale jets are still debated (Harris \& Krawczynski 2006; Zhang et al. 2010; Meyer et al. 2015; Zargaryan et al. 2017). 4C +49.22 is a $\gamma$-ray flat spectrum radio quasar (FSRQ) at redshift $z=0.334$ (Burbidge 1968; Lynds \& Wills 1968). It has a one-side, knotty and wiggling jet, and its knots were resolved at the radio, optical and X-ray bands (Owen \& Puschell 1984; Akujor \& Garrington 1991; Sambruna et al. 2004, 2006). This source was not detected by the previous $\gamma$-ray detectors, such as EGRET (Hartman et al. 1999) and AGILE (Pittori et al. 2009), but a bright $\gamma$-ray outburst was detected with \emph{Fermi}/LAT (Reyes et al. 2011; Cutini et al. 2014). The outburst was also simultaneously observed from the microwave to the X-ray bands with \emph{Planck} and \emph{Swift}. The $\gamma$-ray flux is highly variable and correlated with the emission in the low energy bands, indicating that the $\gamma$-ray outburst is from the compact core region (Cutini et al. 2014). In addition, a new component from the radio-core around the time of the $\gamma$-ray outburst was catched with the Very Long Baseline Array (VLBA, Cutini et al. 2014). This robustly suggests that the outburst is in the vicinity of the core region and is related to the activities of the central black hole. Interestingly, a steady $\gamma$-ray emission component was detected with the \emph{Fermi}/LAT during the past 8 operation years. It is unclear whether the steady $\gamma$-ray emission component is attributed to the radio-core or the knots of this source. So far, the $\gamma$-ray emission outside the radio-core was only convincingly detected by the \emph{Fermi}/LAT for the radio lobes of Cen A (Abdo et al. 2010) and Fornax A (McKinley et al. 2015; Ackermann et al. 2016). If the steady $\gamma$-ray emission of 4C + 49.22 is from the knots, it wound be added as a valuable source with $\gamma$-ray emission at the large-scale jet structure. This paper dedicates to study the emission mechanisms of the radio-core and knots for 4C +49.22 for revealing the jet properties from the radio-core to large-scale knots. We analyzed the observational data of \emph{Fermi}/LAT for 4C +49.22 in the past 8 years, and the derived $\gamma$-ray light curve is presented in Section 2. We model the broadband SEDs of the core region at different epochs with a single-zone leptonic model in the IC/BLR scenario (Section 3.1). We also model the SEDs in the radio--optical--X-ray band for the knots with the leptonic model in the IC/CMB scenario and compare the $\gamma$-ray flux predicted by the model to the steady $\gamma$-ray component of the LAT observation (Section 3.2). The jet properties from the core region to the large-scale knots are presented in Section 4. Discussion and a summary are given in Section 5 and Section 6, respectively. Throughout, $H_0=71$ km s$^{-1}$ Mpc$^{-1}$, $\Omega_{\rm m}=0.27$, and $\Omega_{\Lambda}=0.73$ are adopted. | We dealt with and analyzed the long-term monitoring data of 4C +49.22 by \emph{Fermi}/LAT, besides a large outburst, and a long-term, steady $\gamma$-ray emission component is observed in the $\gamma$-ray light curve, which almost can be fitted by a constant flux. The broadband SEDs of the core region at different $\gamma$-ray emission epochs (during the large flare, post-flare, low state) can be well reproduced with the single-zone leptonic model, synchrotron+SSC+IC/BLR. The SEDs of the eight knots in large-scale can also be well explained with single-zone leptonic model, synchrotron+SSC+IC/CMB, and the $\gamma$-ray fluxes predicted by the model are lower than that of the steady $\gamma$-ray emission observed with \emph{Fermi}/LAT. The synthetical fluxes predicted by the model for the eight knots at the \emph{Fermi}/LAT energy band are still lower than this steady $\gamma$-ray emission with the integral flux ratio of 0.21. It indicates that the steady $\gamma$-ray emission is still dominated by the core radiation, but may be partially contributed by the large-scale knots, which may conceal the low-level flux variation of the $\gamma$-ray emission from the core region. The decreases of the flux at X-ray and optical bands and the derived bulk Lorentz factors along the jet in large-scale may indicate that the jet decelerates in large-scale by interacting with surrounding medium. On the basis of the fitting parameters, we calculated the jet powers and the power carried by each component in pc-scale and kpc-scale. It was found that the jet powers independently estimated for pc- and kpc-scale jets of 4C +49.22 are roughly in the same order. The magnetization parameters and radiation efficiencies of the core region decrease with the decrease of the $\gamma$-ray emission flux, and the jet changes from Poynting flux dominated into particle dominated. Comparing with the core region, the knots in large-scale are particle dominated with very low radiation efficiencies. Therefore, the Poyting flux dominating jet in pc-scale converts to the particle dominating jet in kpc-scale through some unknown mechanisms. The synchrotron+SSC+IC/CMB model indeed induces the `super-Eddington" jet powers for some sources as suggested by some authors (e.g., Dermer \& Atoyan 2004; Uchiyama et al. 2006; Meyer et al. 2015). | 18 | 8 | 1808.07978 |
1808 | 1808.02593_arXiv.txt | The tilt, rotation, or offset of each CCD with respect to the focal plane, as well as the distortion of the focal plane itself, cause shape distortions to the observed objects, an effect typically known as field distortion (FD). We point out that FD provides a unique way of quantifying the accuracy of cosmic shear measurement. The idea is to stack the shear estimators from galaxies that share similar FD-induced shape distortions. Given that the latter can be calculated with parameters from astrometric calibrations, the accuracy of the shear estimator can be directly tested on real images. It provides a way to calibrate the multiplicative and additive shear recovery biases within the scientific data itself, without requiring simulations or any external data sets. We use the CFHTLenS images to test the Fourier\_Quad shear recovery method. We highlight some details in our image processing pipeline, including background removal, source identification and deblending, astrometric calibration, star selection for PSF reconstruction, noise reduction, etc.. We show that in the shear ranges of $-0.005\lsim g_1\lsim 0.005$ and $-0.008\lsim g_2\lsim 0.008$, the multiplicative biases are at the level of $\lsim 0.04$. Slight additive biases on the order of $\sim 5\times 10^{-4}$ ($6 \sigma$) are identified for sources provided by the official CFHTLenS catalog (not using its shear catalog), but are minor ($4 \sigma$) for source catalog generated by our Fourier\_Quad pipeline. | \label{intro} Recent weak lensing measurements have indicated marginal tensions with results from the cosmic microwave background \citep{hildebrandt2017,kohlinger2017,DES1,DES2,HSC1}, and galaxy formation models \citep{leauthaud2017}. It would be extremely exciting if these discrepancies point to new physics beyond the concordance $\Lambda$CDM cosmological model, e.g., dynamical dark energy or modified gravity theories \citep{amendola2018}. There are, however, also concerns that weak lensing measurements still contain systematic errors that have not been identified and corrected, regarding either shape measurement or photo-z errors \citep{EL2018}. To establish weak lensing as a robust cosmological probe, ways of testing the accuracy and consistency of the measurements are indispensable. Simulations are typically used to calibrate the shear recovery accuracy \citep{great3}, but image processing at different stages could all introduce problems that are not easily identified, and therefore not included in the simulations. There are also ways of testing the consistency of shear recovery based on real data. For example: cross-correlation between the PSF shape and the galaxy shape is used to find out if there are residual anisotropic PSF effects that are not removed \citep{kaiser1995,fischer2000}; E/B mode decomposition is used to check if the lensing signal has a gravitational origin \citep{crittenden2002,schneider2002}. Nevertheless, we find that these tests are not very sensitive to the multiplicative biases that people typically require to correct for the sensitivity of their shear estimator to the underlying shear signal. Given that the multiplicative biases are directly degenerate with the amplitudes of the lensing signals and the shear-shear correlations, and thereby several key cosmological parameters such as $\sigma_8$ and $\Omega_m$, it is desirable to calibrate the multiplicative biases based on real data as an alternative to the existing methods \citep{vallinotto2011,zhangpengjie2015}. The purpose of this work is to propose a solution with the help of field distortion, an effect that exists universally in optical systems. The effect of field distortion modifies the galaxy image in a way similar to that of lensing. The amplitude and direction of this distortion can be calculated using the parameters derived from astrometric calibration on one hand, and recovered from the galaxy images on the other hand. We therefore have a natural way of testing the accuracy of shear recovery directly using real data, without requiring simulations or external data sets. In \S\ref{fd}, we introduce the concept of field distortion, and the way of deriving it using the astrometric parameters. In \S\ref{shear_measure}, we give a brief introduction of the Fourier\_Quad shear measurement method \citep{zlf2015}. We use the CFHTLenS data \citep{heymans2012,erben2013} to demonstrate our idea. The details regarding our image processing of the CFHTLenS data are given in \S\ref{reduction}. It highlights a few key steps in our pipeline. Our main results are shown in \S\ref{results}. Finally, in \S\ref{conclusion}, we give brief conclusions as well as discussions regarding the current status of our pipeline and future prospects. | \label{conclusion} We have presented a way of testing shear recovery accuracy on real images. The target "shear signals" are given by the field distortion effect, which is naturally involved in any optical system. The distribution of the FDS can be accurately mapped out on the CCD plane with the astrometric parameters. These shear signals are recoverable with galaxy images whenever there are a large number of exposures available. The cosmological shear signals in this case are generally not relevant. By comparing the FDS with the stacked galaxy shear estimators, one can directly quantify the multiplicative and additive shear biases on real images. As an example, we show the performance of the Fourier\_Quad method using the CFHTLenS data. The FDS values on MegaCam are around $0.005$ or less, very suitable for shear calibration at the level of cosmic shear. We have described in \S\ref{reduction} a number of important details in the image processing pipeline of Fourier\_Quad, including background removal, source identification and deblending, astrometric calibration, star selection for PSF reconstruction, noise reduction. Our pipeline can also read in source locations (RA \& DEC) from an external source catalog. Our main results are presented in fig.\ref{result_FQ} and fig.\ref{result_LF} using two different source catalogs: the FQ one is by our own pipeline, and the LF catalog is from the official data release of CFHTLenS. The shear signals in both cases are measured by Fourier\_Quad. Overall, the multiplicative biases are at the level of $\lsim 4\%$. The results from the sources of the LF catalog reveal some slight additive biases of about $5\times10^{-4}$ ($6\sigma$ significance) for both shear components. This problem is minor for the FQ catalog: only $g_2$ has an additive bias of $3.2\times 10^{-4}$ ($4\sigma$). We believe the difference is due to some subtleties in the source locating/identification part of the pipeline of CFHTLenS, which is very different from ours. The most significant difference is that sources in the LF catalog are identified on co-added images, while our pipeline so far processes every exposure individually. A combinatory multi-exposure version of Fourier\_Quad is still under development. Within the LF catalog provided by CFHTLenS, we find that there are a large portion of sources that have valid magnitudes, but not shear measurements. This problem is partially remedied by our Fourier\_Quad pipeline, which provides valid shear measurements for most of the sources in the LF catalog, as shown in fig.\ref{comp_w2m0m0}. In principle, for each source, Fourier\_Quad can yield one shear measurement from each exposure. The reality is that some sources in the LF catalog cannot be properly identified on single exposures in Fourier\_Quad, especially at the faint end as shown in fig.\ref{w2m0m0_expos}. This is due to their weak significance or image defects. For valid source images, their shear estimators can all be measured by Fourier\_Quad. The only exception is when there are not enough bright stars on the chip for PSF reconstruction. The shear recovery part of our pipeline simply does not apply any selection rules on the source morphology (shape, size, etc.), allowing us to avoid complicated decisions regarding selection effects due to shear measurement itself. Note that it is not even necessary to exclude stellar objects or point sources from the galaxy samples in Fourier\_Quad. We consider this a significant advantage of our pipeline. Our proposal of shear testing with field distortion should be useful for any shear measurement algorithm, including some recently developed ones \citep{sh2017,hm2017,of2017,tewes2018,pujol2018,li2018}. It is better to work with individual exposures, in which case the field distortion signals are well defined. For algorithms like Lensfit, which simultaneously fits the source shape on multiple exposures, the field-distortion signals are already included in the forward modeling of the galaxy image on each exposure. In this case, the recovered ellipticities can still be plotted against the average field-distortion signal to form a null test, assuming the spatial offsets of the relevant exposures on the sky are much smaller than their sizes. Further discussion of these topics is beyond the scope of this work. It is also important to note that accurate astrometric calibration is indispensable for a successful shear measurement program. We have introduced a modification to the fitting between the projected sky plane and the CCD plane in \S\ref{astrometric_calibration}, enabling us to easily extend the fitting functions to high order polynomials, and to check the convergence of the astrometric solution. A more comprehensive comparison between the astrometry part of our pipeline and the standard software, such as SCAMP, will be provided in a separate work. We also caution that polynomial fitting is not likely good enough for recovering distortions caused by instrumental effects at special locations \citep{bernstein2017} of the CCD. For the CFHTLenS data, we have not studied such a problem carefully. We leave a detailed discussion of this topic to a future work. Finally, it is encouraging to note that our Fourier\_Quad pipeline is now able to carry out shear measurement from almost raw CCD images (after flat-field correction). The overall image processing speed is very fast. For example, for the CFHTLenS data, overall, it only takes about 0.02 CPU*seconds for each galaxy image on average. We therefore consider the Fourier\_Quad pipeline a very promising tool for probing the cosmic structure in the ongoing and planned large scale galaxy surveys. In the future development of our pipeline, we will try to include the following components: 1. source identification using information from multiple exposures; 2. ways of stacking under-sampled images for accurate shear recovery (mostly for space-based missions); 3. treatments of instrumental effects \citep{rhodes2010,antilogus2014}. | 18 | 8 | 1808.02593 |
1808 | 1808.10453_arXiv.txt | The gas-phase metallicity of low-mass galaxies increases with increasing stellar mass ($M_\ast$) and is nearly constant for high-mass galaxies. Theory suggests that this tight mass-metallicity relationship is shaped by galactic outflows removing metal-enriched gas from galaxies. Here, we observationally model the outflow metallicities {of the warm outflowing phase from} a sample of seven local star-forming galaxies {with stellar masses between 10$^{7}$--10$^{11}$~M$_\odot$}. We estimate the outflow metallicities using four weak rest-frame ultraviolet absorption lines, the observed stellar continua, and photoionization models. The outflow metallicity is flat with $M_\ast$, with a median metallicity of $1.0\pm0.6$~Z$_\odot$. The observed outflows are metal-enriched: low and high-mass galaxies have outflow metallicities 10-50 and 2.6 times larger than their ISM metallicities, respectively. {The observed outflows are mainly composed of entrained ISM gas {with} at most 22\% of the metals directly coming from recent supernovae enrichment}. The metal outflow rate shallowly increases with $M_\ast$, as $M_\ast^{0.2 \pm 0.1}$, because the mass outflow rate shallow increases with $M_\ast$. Finally, we normalize the metal outflow rate by the rate at which star formation retains metals to calculate the metal-loading factor. The metal-loading factor inversely scales with $M_\ast$. The normalization and scaling of the metal-loading factor agree with analytic expressions that reproduce observed mass-metallicity relations. Galactic outflows fundamentally shape the observed mass-metallicity relationship. | Many galaxy properties are highly correlated. The star formation rate (SFR) sublinearly increases with the stellar mass (\mstarp), at nearly all redshifts, as the star formation main sequence \citep{brinchmann2004, salim07, noeske07, elbaz07}; the stellar mass increases with the halo mass \citep{moster10}; the gas-phase metallicity scales with \mstar \citep{Lequeux, skillman, tremonti04, erb06, berg, andrews13, zahid}; and the gas-phase metallicity correlates strongly with \mstar {\it and} SFR, as the so-called fundamental metallicity relation \citep{mannucci, lopez10}. By connecting different galactic components, these relations suggest that physical processes shape the formation and evolution of galaxies. \\ Among these relations, the stellar mass gas-phase metallicity relation (MZR) has perhaps received the most attention. Observations suggest that the gas-phase metallicity (typically measured from strong optical oxygen nebular emission lines as 12+log(O/H)) increases with the \mstar of low-mass galaxies, but flattens to a nearly constant value at large \mstar \citep{tremonti04}. While this qualitative shape is nearly universally observed, the normalization (i.e. the absolute metallicity values and the \mstar where the relations flatten) strongly depends on how the metallicity is calculated \citep{kewley08}. Absolute metallicity calibrations can vary up to 0.7~dex for different calibration systems. \\ The MZR connects host galaxy observable properties to fundamental processes in galaxy evolution. To start, the quantities relate the integrated conversion of gas into stars (the total \mstarp) to the integrated amount of gas processed through, and therefore enriched by, stars (the metallicity). While both \mstar and metallicity are by-products of star formation, they represent different phases of baryonic matter. The simplest attempt to theoretically explain the MZR is the closed box scenario where a galaxy initial consists of metal-poor gas which gravitationally collapses to form stars \citep[see the review in ][]{tinsley80}. No gas enters or leaves the galaxy in this scenario. In this simplistic situation, high-mass stars produce metals through nucleosynthesis, returning those metals, at a yield of $y$, to the surrounding gas through stellar winds and supernovae. With these assumptions, the gas-phase metallicity (\zism) increases as gas is processed through stars as \begin{equation} Z_\mathrm{ISM} = y \ln\frac{M_g + M_\ast}{M_g} = y\ln\mu^{-1} , \label{eq:closed} \end{equation} where M$_g$ is the gas mass and $\mu$ is the gas mass fraction \citep{searle}. While simple, \autoref{eq:closed} catastrophically fails to reproduce the metal content of stars in the solar neighborhood \citep[the G-dwarf problem;][]{schmidt, vandenbergh, pagel75}; other physics must shape the metallicities of galaxies. Galaxies do not evolve as closed systems. Relatively metal poor gas accretes onto galaxies from the intergalactic medium (IGM), diluting the metallicity \citep{larson72}. Accretion rates larger than the SFR may decrease the metallicity of a galaxy to levels observed in low-mass galaxies \citep{edmunds, koppen, garnett}. However, these excessive accretion events lead to bursts of star formation that enrich the galaxy back to the closed box value (\autoref{eq:closed}). Accretion cannot decrease metallicities over an extended duration \citep{dalcanton}. While galaxies acquire gas and metals from the IGM, they also lose metals through galactic outflows \citep{heckman90, heckman2000, veilleux05}. Supernovae, high-energy photons, cosmic rays, and stellar winds inject energy and momentum into the ISM, accelerating gas out of galaxies. Since low-mass galaxies have shallower gravitational potentials, star formation powered galactic outflows more efficiently remove metals from low-mass galaxies to produce the MZR \citep{larson74, dekel86, tremonti04, dalcanton, finlator08, peeples11, lilly13}. In fact, analytic work by \citet{dalcanton} showed that galactic outflows are the most likely physical process to decrease the metallicites of low-mass galaxies. To produce the MZR, galactic outflows must remove more metals than their star formation retains \citep{dalcanton, peeples11}. By preferentially driving metals out of galaxies, stellar feedback may decrease the gas-phase metallicity of low-mass galaxies and shape the MZR. Galactic outflows are ubiquitously observed in star-forming galaxies at all epochs as broad blue shifted emission lines \citep{lynds, sharp, arribas2014} and interstellar metal absorption lines \citep{heckman90, heckman2000, pettini2002, shapley03, rupke2005b, martin2005, weiner, chen10, martin12, rubin13, bordoloi, heckman15}. However, it is challenging to measure the amount of metals and mass removed by galactic outflows (the mass outflow rate; \moutp) because outflows have uncertain geometries, ionization corrections, and metallicities; all of which may add an order of magnitude uncertainty to \mout estimates \citep{murray07, chisholm16b}. Observations need to determine if the metal outflow rate (\mzp) of galactic outflows is sufficient to shape the MZR. In a series of papers, we have explored galactic outflows from local star-forming galaxies using restframe ultraviolet (UV) down-the-barrel spectroscopy from the Cosmic Origins Spectrograph (COS) on the {\it Hubble} Space Telescope. We started this series by characterizing the host galaxies, the stellar continua, and the galactic outflow properties using a single \siii absorption line \citep[][ hereafter Paper I]{chisholm15} and multiple absorption lines spanning a range of ionization states \citep[][ hereafter Paper II]{chisholm16}. In \citet{chisholm16b} (hereafter Paper III) we used the observed column densities to map out the ionization structure, outflow metallicity, and \mout of a single galaxy, NGC~6090. \citet{chisholm17} (hereafter paper IV) extended the \mout calculation to a sample of 7 galaxies, with spectra that have sufficient signal-to-noise ratios to explore how \mout scales with \mstarp. \begin{table*} \caption{{Derived galaxy properties. Column 1 is the galaxy name; column 2 is the stellar mass (\mstarp); column 3 is the galactic circular velocity (\vcircp) calculated using a Tully-Fischer relation \citep{reyes}; column 4 is the star formation rate of the entire galaxy (SFR); column 5 is the SFR within the COS aperture (SFR$_\mathrm{COS}$); column 6 is the oxygen abundance (12+log(O/H)); column 7 is the stellar continuum metallicity (Z$_\mathrm{s}$); column 8 is the light-weighted stellar age; column 9 is the stellar continuum attenuation (E$_\mathrm{S}$(B-V)) corrected for foreground reddening \citep{schlegel}; column 10 {gives the distances used and, when appropriate, the references for the distances}; column 11 lists the \textit{HST} proposal ID for each observation; and in column 12 we give references to other papers that use the COS data. Columns 7-9 were determined from the stellar continuum fits in \autoref{cont}.Metallicity references: a) \citet{cortijo}, b) \citet{izotov97}, c) \citet{izotov98} d) \citet{james13}, e) \citet{ostlin}, f) \citet{perez}.} {Distance references in column 10 are: ($\alpha$) \citet{aloisi05} and ($\beta$) \citet{fiorentino}; the rest are computed from the redshift and the Hubble flow.} References in column 12 are: (1) \citet{alexandroff}, (2) \citet{duval}, (3) \citet{france2010}, (4) \citet{fox2013}, (5) \citet{fox14}, (6) \citet{hayes14}, (7) \citet{heckman15}, (8) \citet{james}, (9) \citet{claus2012}, (10) \citet{ostlin}, (11) \citet{pardy}, (12) \citet{Rivera}, (13) \citet{richter2013}, and (14) \citet{wofford2013}.} \resizebox{\textwidth}{!}{ \begin{tabular}{lccccccccccc} \hline Galaxy name & log(\mstarp) & \vcircp & SFR & SFR$_\mathrm{COS}$ & 12+log(O/H) & Z$_\mathrm{s}$ & <Age> & E$_\mathrm{S}$(B-V) & D & Proposal ID & References \\ & [log(M$_\odot$)] & [\kmsp] & [M$_\odot$~yr$^{-1}$]& [M$_\odot$~yr$^{-1}$] & [dex] & [Z$_\odot$] &[Myr] & [mag]& [Mpc] & & \\ (1) & (2) & (3)& (4)& (5)& (6) & (7) &(8) & (9) & (10) & (11) & (12) \\ \hline SBS~1415+437 &6.9 & 18& 0.016 & 0.016 & 7.59$^c$ & 0.05 & 3.7 & 0.10 &13.6$^{\alpha}$& 11579 &8 \\ 1~Zw~18 &7.2 & 21 & 0.023 & 0.023 & 7.22$^b$ & 0.05 & 7.3 & 0.09 &18.2$^{\beta}$ & 11579 & 8 \\ MRK~1486 &9.3 & 82 &3.6 & 2.5& 7.80$^e$ & 0.2 &8.0 & 0.22 & 148& 12583 & 2, 6, 10, 11, 12 \\ KISSR~1578 & 9.5 & 94 & 3.7 & 2.1 & 8.07$^e$ & 0.2 & 5.5 & 0.13 & 122& 11522 & 3, 14 \\ Haro~11 & 10.1 & 137 & 26 & 12 & 7.80$^d$ & 0.4 & 10.1 &0.13 & 89&13017& 1, 7 \\ NGC~7714 &10.3 & 156 & 9.2 & 3.1 & 8.23 $^f$ &1 & 4.3 & 0.31 & 40& 12604 & 4, 5, 13\\ NGC~6090 & 10.7 & 202& 25& 5.5 & 8.40$^a$& 1 & 4.5 & 0.32& 128& 12173 &9\\ \end{tabular} } \label{tab:sample} \end{table*} Here, we conclude the series of papers by examining, for the first time, the outflow metallicities and {\it metal} outflow rates of star-forming galaxies. We begin by reviewing the host galaxy properties and the COS observations of our sample (\autoref{data}). We then summarize the steps to determine the metal outflow rates of these seven galaxies, including: the measurement and removal of the stellar continuum (\autoref{cont}), the profile fitting and mass outflow rate calculation (\autoref{proffit}), and the ionization modeling (\autoref{cloudy}). Trends with our metal outflow rates are subsequently explored (\autoref{results}). We then compare our estimates to previous observations (\autoref{prev}), discuss the role outflows play in shaping the observed MZR (\autoref{mzrel}), and explore some limitations and future steps of this study (\autoref{further}). In \autoref{cloudy_obs}, we discuss the full parameter space of our photoionization models. These observations provide the first empirical evidence that galactic outflows shape the mass-metallicity relationship. \\\ Throughout this paper we use a solar gas-phase metallicity of 0.0142 and 12+log(O/H)$_\odot$ of 8.69 \citep{asplund}. Stellar solar metallicity is 0.02 \citep{claus99}. {Additionally, we use the Chabrier initial mass function \citep[IMF;][]{chabrier} when determining host galaxy properties and a Kroupa IMF \citep{kroupa} when fitting the stellar continua in \autoref{cont}.} | \subsection{Comparison to previous results and simulations} \label{prev} \subsubsection{Previous observations} The local galaxies in our sample are well-studied, and previous studies have estimated similar properties as our photoionization models. For instance, log($N_\mathrm{HI}[\mathrm{cm}^{-2}]$) was estimated for MRK~1486 \citep{duval} and Haro~11 \citep{rivera-thorsen17} by modeling the \siii absorption lines and the \lya\ emission profiles. These previous studies found log($N_\mathrm{HI}[\mathrm{cm}^{-2}]$) of 19.5 and $> 19.0 \pm 0.2$, respectively. These independent values agree with the estimated log(\nhi) found from our {\small CLOUDY} modeling (see \autoref{tab:ion}). Other studies carefully measured the metallicity from the low metallicity galaxy 1~Zw~18 using the observed \lya\ profile and the sulphur absorption lines to determine a metallicity of $\sim0.02$~Z$_\odot$ in the ionized gas and $\sim0.03$~Z$_\odot$ in the neutral gas \citep{lebouteiller}. This metallicity is substantially lower than the $1.6\pm0.3$~Z$_\odot$ that we derived using the {\small CLOUDY} models, and similar to \zism. Similarly, \citet{james} estimated log(\nhi)$ = 21.1$ from SBS~1415+437. The \ion{H}{i} column densities and metallicities measured here are significantly different from these studies. The \lya\ regions for 1~Zw~18 and SBS~1415+437 are extremely difficult to interpret due to the presence of damped Milky Way absorption \citep{kunth}, geocoronal emission, absorption from high-velocity clouds \citep{hvc1, hvc2}, \lya\ plus nebular emission from the background galaxy \citep{lebouteiller}, and stellar continuum features. Importantly, the broad \ion{N}{v}~1240\AA\ P-Cygni feature and stellar \ion{H}{i} absorption can be interpreted as damped Ly$\alpha$ wings. \citet{james} and \citet{lebouteiller} carefully account for many of these features, but \citet{james} do not fit the stellar continuum and \citet{lebouteiller} fit the stellar continuum after measuring the \ion{H}{i} column density. Not accounting for the stellar continuum may over-estimate the \ion{H}{i} column density and under-estimate the metallicity. Additionally, we accounted for partial covering when determining the metal column densities while \citet{lebouteiller} and \citet{james} assumed a unity covering fraction. We inferred the covering fractions from the observed absorption profiles, thus the covering fractions should be accounted for when calculating column densities. Since both 1~Zw~18 and SBS~1415+437 have small \siiv covering fractions (see \autoref{tab:prof}), assuming the lines are fully covered systematically decreases the estimated column densities from Voigt fits. This may have led to the large differences between the measured metallicities. Further, the photoionization models of \citet{lebouteiller} are only observationally constrained by the low-ions (up to \siiip). The authors find a factor of five difference in the metallicities derived in the neutral and ionized Si zones. By not observationally constraining the ionization models with high-ions, it is possible that the ionization corrections of the high-ions could be even larger. Our ionization models suggest that 58\% of the Si from 1~Zw~18 is in the \siiii or \siiv transitions; not including higher ionization phases may under-predict the total amount of outflowing metals. {\small CLOUDY} models with 0.01~Z$_\odot$ under-predict the \siiv column densities by over an order of magnitude for all U and \no\ combinations, even though they allow for the large log(\nhi) values found by previous studies (the {\small CLOUDY} models exceed log(\nhi) = 26). In \citetalias{chisholm16}, we also found that very low \zo were incompatible with the observed ionization structure of the outflows (fig.~14 of that paper). Further, the low metallicity {\small CLOUDY} models produce larger \siivp/\siii column density ratios than are observed (see \autoref{cloudy_obs} and \autoref{fig:cloudy_mod} below). Another possibility is that these two low-mass galaxies have substantial amounts of low-metallicity gas at zero-velocity, and the strong damped \lya\ absorption traces this systemic gas which has the ISM metallicity. The metal absorption lines are from an enriched galactic outflow, while the \lya\ from the outflow is blended with the systemic damped \lya. The deepest portion of the low-ionization metal absorption lines from 1~Zw~18 are blueshifted, relative to the 21~cm line, by 15~\kms \citep{lebouteiller}, while the equivalent width weighted velocity of the \siivp~1402\AA\ line is blueshifted by $-169\pm15$~\kms (see \autoref{fig:s42}). The low-velocity \hi\ absorption is easily blended with zero-velocity gas, making it challenging to disentangle the relatively weak outflow--as traced by the metal lines--from the strong systemic \lya\ absorption. Zero-velocity contamination is a significant problem for outflow studies \citep{weiner, chen10, rubin13}, and it could play a large role in the discrepancies between the metallicities measured here and previous studies. To this end, \citet{lelli} observed 21~cm emission from 1~Zw~18 with the VLA. These authors modeled the \ion{H}{i} emission maps as a disk plus a radial flow with a best fit velocity of 15~\kms \citep[the same as the low-ionization absorption lines observed by][]{lebouteiller}. Interestingly, the \hi\ disk is not centered on the star-forming region probed by the COS aperture, possibly suggesting that the \hi\ and star formation are not coincident. Once they removed the disk component, \citet{lelli} found a residual log(\nhi[cm$^{-2}$])~$=19.6$, corresponding to outflowing gas (their fig. 8), consistent with the ionization modeling here. Therefore, it is plausible that the metal absorption lines from the two low-mass galaxies trace a metal-enriched outflow, while the \lya\ absorption traces a large-scale, low-metallicity, zero-velocity disk. Other studies have measured outflow metallicities, but not in the warm photoionized outflow phase that we observed. \citet{martin02} fit the integrated X-ray spectrum of a low-metallicity (\zism~$=0.2$~Z$_\odot$; 12+log(O/H) = 8.0; log(\mstarp)~$=7.8$) galaxy, NGC~1569, with a four component model: a disk, a power-law, an absorbed thermal component, and an unabsorbed thermal component. Both thermal components were treated as the outflow. The authors minimized the fit to the X-ray spectrum when the $\alpha$ elements (like Si and O used here) had \zop~$=1$~Z$_\odot$, {but \zo is not tightly constrained due to a broad \zo distribution (see their fig. 19). Regardless, the solar metallicity measurement of the \textit{hot} outflow is} in agreement with the median outflow metallicity we measured {for the warm outflow phase} in \autoref{results}, and 5 times larger than \zism. X-ray studies of 10$^7$~K plasma from the local starbursting galaxy M~82 find super-solar metallicities that are >$3$~Z$_\odot$ \citep{strickland09}. This large \zo likely arises because the large-scale X-ray emission traces pure supernovae ejecta, whereas the COS metal absorption lines trace a mixture of ISM gas and supernovae ejecta. Removing ISM gas drives the MZR, thus tracing the metallicity of the gas being removed is crucial for determining which physical process creates the MZR. We return to the idea of mixing the two phases in \autoref{highz}. \subsubsection{Simulated metal outflow rates} \label{sim} The \mz measures how rapidly outflows remove metals from galaxies. Simulations usually tune \mz to match the observed MZR (see the discussion in \autoref{mzrel}). However, recent simulations from the Feedback In Realistic Environments simulations \citep[FIRE;][]{hopkins14} enrich gas according to the supernovae yields \citep{woosley95}, while driving feedback by injecting energy and momentum into the gas according to stellar population synthesis models. The metal production and feedback processes successfully reproduce the MZR at $z = 0-3$ \citep{ma}. \citet{muratov17} measure \mz at 0.25~R$_\mathrm{vir}$ for a small sample of these simulated galaxies. Two galaxies in particular compare well with our sample because they have \mstar of $2\times10^9$ and $6\times 10^{10}$~M$_\odot$ at $z = 0$ (their m11 and m12i simulations; {the values come from the peaks in the left panel of their fig.~1}). These galaxies have \mz peak values between $0.01-0.04$~\sfr for the low-mass galaxy, and $0.05-0.11$~\sfr for the high-mass galaxy (after converting to the solar metallicity conventions used here). These two simulated galaxies are comparable to MRK~1486 (comparable to m11, the low-mass FIRE galaxy) and NGC~6090 (comparable to m12i, the high-mass FIRE galaxy), which have \mz of $0.02\pm0.01$ and $0.06\pm0.02$~\sfrp, respectively. {The low-mass FIRE galaxy, m10, has \mstar at $z = 0$ three times less massive than the lowest mass galaxy in our sample (2.3$\times10^{6}$~M$_\odot$), which makes it challenging to directly compare to our sample. {However, if we extrapolate} the fitted relationship between \mz and \mstar in \autoref{eq:metaloutflow}, we predict a \mz of 0.0031$^{+0.0032}_{-0.0016}$~\sfr for a galaxy with \mstar of 2.3$\times10^{6}$~M$_\odot$, consistent, within the large errors, with the maximum of the \mz range of 0.0002-0.0016~\sfr for m10.} The simulated \mz values are surprisingly consistent with our observations. In \autoref{mzrel} we discuss the implications of the \mz values and why they determine whether simulations reproduce the observed MZR. \subsection{Shaping the mass-metallicity relation with galactic outflows} \label{mzrel} \begin{figure*} \includegraphics[width = \textwidth]{mz-relation.pdf} \caption{Relationships that shape the mass-metallicity relation. {\it Left panel:} Outflow metallicity (\zop) normalized by the gas-phase metallicity of the galaxy (\zism) versus the stellar mass of the galaxy (\mstarp). {\it Right panel:} The logarithm of the metal-loading factor, the ratio of metals removed by the galactic outflow to metals locked into stars, versus the logarithm of \mstarp. The metal-loading factor inversely scales with \mstar at the 3$\sigma$ significance level (p-value < 0.001). Included in the right panel are two analytic predictions of the metal-loading factor required to match the observed MZR from \citet{denicolo} (red dashed line) and \citet{tremonti04} (blue dot-dashed line), as fit by \citet{peeples11}. The observed metal-loading factors closely follow the analytic relation required to reproduce the observed \citet{denicolo} MZR. } \label{fig:zout} \end{figure*} Here we explore the implications that the observed metal outflows have on the stellar mass gas-phase metallicity relation (MZR). The left panel of \autoref{fig:z} indicates that \zo does not vary with \mstarp, rather \zo scatters about a median of 1.0~Z$_\odot$. However, when \zo is normalized by the observed ISM gas-phase metallicity (Z$_\mathrm{ISM}$; left panel of \autoref{fig:zout}), the outflow metallicities from six of the seven galaxies are more metal-enriched than their ISM metallicities. Even at {the upper-end of the mass range of our sample} (log(\mstarp[M$_\odot$)~$>9$), the median \zo is 2.6 times larger than \zism. The outflows of the lowest mass galaxies carry out even more metals as their outflow metallicities are 12 and 46 times larger than their ISM metallicities. The observed outflows are metal-enriched. \citet{dalcanton} analytically studied three possible origins of the MZR: inflows of low-metallicity gas, low-metallicity outflows, and metal-enriched outflows. The author finds that outflows with \zop/\zism~>~1 are the only way to decrease the metallicity of low-mass galaxies to the observed levels \citep[a similar conclusion was found by][]{lu}. However, \citet{dalcanton} did not consider the possibility that outflows remove more mass in lower mass galaxies than in higher mass galaxies. \citet{peeples11} introduced the metal-loading factor, which describes how efficiently galactic outflows remove metals relative to how efficiently star formation retains metals. Numerically, this is given as \begin{equation} \zeta = \frac{Z_\mathrm{o}}{Z_\mathrm{ISM}} \frac{\dot{\mathrm{M}}_\mathrm{o}}{\mathrm{SFR_\text{COS}}} , \label{eq:zeta} \end{equation} where we used \sfrc\ in \autoref{eq:zeta} because we assumed that the local star formation energy and momentum accelerates the outflow (see \autoref{sample}). The numerator in \autoref{eq:zeta} is the rate at which outflows remove metals and the denominator is the rate at which star formation retains metals in new stars. In essence, what drives the MZR is the competition between the ejection and retention of metals. This tension is the metal-loading factor (right panel of \autoref{fig:zout}). The metal-loading factor strongly anti-correlates with \mstar (3$\sigma$, p-value < 0.001, relation; Pearson's $r$ of -0.97). To better understand this competition, consider how the metallicity of a galaxy changes with time \citep{larson74, tinsley80, dalcanton, finlator08, peeples11, lilly13, zahid, lu}. Galaxies accrete low-metallicity gas from the IGM (with a metallicity of Z$_\mathrm{IGM}$) at a rate $\dot{M}_\mathrm{acc}$. This low-metallicity gas collects in the gravitational potential of the galaxy and forms stars. These stars produce metals (at a yield $y$) and instantly return a fraction, $R$, of these metals to the ISM through stellar winds or supernovae. The same stellar winds and supernovae inject energy and momentum into the newly enriched gas to drive a galactic outflow which removes a mass of metals at a rate of $\dot{M}_z= Z_\mathrm{o}\times\dot{M}_\mathrm{o}$. This approach, often called the leaky box (or bathtub) model, describes the change in \zism\ as a differential equation: \begin{equation} \frac{\mathrm{dZ}_\mathrm{ISM}}{\mathrm{dt}} = -\mathrm{Z}_\mathrm{ISM} \text{SFR} +\mathrm{Z}_\mathrm{IGM} \dot{M}_\mathrm{acc} -Z_o \dot{M}_o + y (1-R) \text{SFR} . \label{eq:mmt} \end{equation} Assuming that $Z_\mathrm{IGM} = 0$, $y$ is constant, and that metals are instantly recycled back to the ISM, \citet{peeples11} solved this differential equation as \begin{equation} \mathrm{Z}_\mathrm{ISM} = \frac{y}{\zeta +\alpha \frac{M_g}{M_\ast} +1} , \label{eq:peep} \end{equation} where $\alpha$ is a constant of order 1 and ${M_g}$/${M_\ast}$ is the gas-to-stellar mass fraction of the galaxy. \autoref{eq:peep} resembles the closed box model of the MZR (\autoref{eq:closed}), but includes the loss of metals by galactic outflows. Assuming that $y$ is constant, \autoref{eq:peep} shows that the MZR is completely determined by the scaling of the metal-loading factor and the gas mass fraction. If galaxies accrete metal rich gas ($Z_\mathrm{IGM} \ne 0$), possibly due to a galactic fountain or an enriched IGM, then the denominator of \autoref{eq:peep} must include the accretion loading factor ($\frac{\mathrm{Z}_\mathrm{IGM}}{\mathrm{Z}_\mathrm{ISM}} \frac{\dot{M}_\mathrm{acc}}{\mathrm{SFR}}$), but we ignore this effect because the IGM is assumed to have a negligible metallicity. The scaling of the metal-loading factor with the circular velocity (\vcircp; a proxy of the stellar mass) is fit with a power-law as \begin{equation} \zeta = \left(\frac{\mathrm{v}_0}{\mathrm{v}_\mathrm{circ}}\right)^{b}+\zeta_0 . \label{eq:peep_zeta} \end{equation} The shape of the metal-loading factor relation has two portions: (1) a low-mass portion below the turn-over mass (v$_0$) where the metal-loading factor rapidly increases with decreasing \mstar as a power-law with exponent $b$, and (2) a high-mass portion, above v$_0$, that has a constant metal-loading factor ($\zeta_0$). We fit \autoref{eq:peep_zeta} to the observed metal-loading factors and find v$_0=91\pm22$~\kmsp, $b=3.4\pm0.7$, and $\zeta_0=0.74\pm0.39$. \citet{peeples11} placed the scaling of eight different MZRs from the literature into \autoref{eq:peep}, along with the observed $M_g/M_\ast$ scaling from three combined samples \citep{mcgaugh05, leroy08, west09, west10}, and fit for the power-law metal-loading factor required to reproduce the observed MZRs. In the right panel of \autoref{fig:zout}, we over-plot their results for the \citet{tremonti04} (blue dot-dashed line) and \citet{denicolo} (red dashed line) relations. \citet{denicolo} calibrated the [\ion{N}{ii}]/H$\alpha$ ratio as a metallicity indicator from a sample of galaxies with a mixture of T$_e$ based oxygen abundances and photoionization model based abundances. Meanwhile, \citet{tremonti04} used a Bayesian approach to estimate metallicities using all of the available strong optical emission lines and a grid of photoionization models. The differences between the two calibration systems--especially at the high-mass end--reflect the large uncertainties of different metallicity calibrations \citep{kewley08}. Other MZRs have similar shapes as the two presented here \citep[see fig. 6 of][]{peeples11}, and the fits to the observed metal-loading factors are broadly consistent with the predictions from \citet{peeples11}. The observed metal-loading factors agree better with the metal-loading factors required to reproduce the \citet{denicolo} relation. This is not surprising because we used T$_e$ based \zism\ values, similar to the \citet{denicolo} calibration. The low-mass (log(\mstarp)~$=7.8$) galaxy from \citet{martin02} also resides on the curves of \autoref{fig:zout}. With a \zop/\zism~$=5$ and an estimated mass-loading factor of 9 {for the hot outflowing component}. A metal-loading factor of 45 (log$\zeta = 1.65$) resides between both the \citet{denicolo} and \citet{tremonti04} curves in \autoref{fig:zout}. This suggests that the metal-loading factor increases monotonically within the \mstar gap in \autoref{fig:zout}, but future observations are required to confirm this (see \autoref{further}). In \autoref{eq:peep}, the declining $M_g/M_\ast$ and the observed metal-loading factors completely describe the shape of the observed MZR. At low-masses (log($M_\ast$[M$_\odot$])$~< 10$), $M_g/M_\ast$ and the metal-loading factors are greater than 1, and both decrease with increasing \mstarp. Therefore, \zism\ increases with increasing \mstarp. However, for log(\mstarp[M$_\odot$])~>~10, $M_g/M_\ast$ is much smaller than 1 and the metal-loading factor flattens to a constant value. This reduces the denominator in \autoref{eq:peep} to $\zeta_0 + 1$, and \zism\ becomes a constant equal to \begin{equation} \text{Z}_\text{flat} =\frac{y}{\zeta_0+1} . \label{eq:zflat} \end{equation} Consequently, the observed Z$_\text{flat}$ from various MZRs provides $y$ estimates. Using Z$_\text{flat}$ of 0.006 and 0.014 from the \citet{denicolo} and \citet{tremonti04} MZRs, we estimate $y$ to be 0.010 and 0.024. These values agree with the range of oxygen yields $y~=~0.007-0.038$ recently calculated by \citet{vincenzo}. This range depends on the IMF, \zs, and the high-mass cut-off. The implied $y$ range of 0.007--0.038 is most similar to the Salpeter and Kroupa oxygen yields \citep{salpeter, kroupa}. Conversely, the Chabrier IMF \citep{chabrier} predicts substantially larger $y$ than implied by the flat portions of the MZR. Since \zop/\zism\ does not scale with \mstarp, the mass-loading factor (\moutp/SFR) alone determines the metal-loading factor. In \citetalias{chisholm17}, we showed that the mass-loading factor decreases with \mstar because the gravitational potential increases more rapidly with \mstar than the momentum supplied by star formation \citepalias[fig.~5 of][]{chisholm17}. This declining efficiency of star formation driven outflows reduces the metal-loading factor and flattens the MZR. The observed physics of driving galactic outflows shapes the stellar mass-metallicity relation. The scatter of the MZR is small \citep[0.1~dex; ][]{tremonti04}. This may imply that the metal-loading factors and gas mass fractions do not strongly vary at a given \mstar (see \autoref{eq:peep}). However, recent \ion{H}{i} observations find a strong correlation between the \ion{H}{i} gas mass and 12+log(O/H) at fixed \mstarp, implying that the scatter in the MZR is tied to the \ion{H}{i} gas mass \citep{brown}. This suggests that the metal-loading factors either have a small scatter, or that the \ion{H}{i} gas mass and metal-loading factors are related. Future observations of both the metal-loading factors and the gas masses are required to break this degeneracy to determine what causes the scatter in the MZR. \subsubsection{Must outflows from low-mass galaxies be metal-enriched to match the MZR?} \label{highz} The lowest mass galaxies in our sample have \zo exceeding their \zism\ by factors of 10-50 (left panel of \autoref{fig:zout}). This starkly contrasts the more massive galaxies in the sample which have Z$_\text{o}~\sim~{\rm 2Z}_\text{ISM}$. Are these highly enriched outflows required to reproduce the observed MZR? If \zo is the only observed property that changes, then the derived metal-loading factors do not change. This counter-intuitive fact arises because the metal-loading factor is a product of \zo times \moutp. Changes in \zo also propagate to \mout by a factor of \zop$^{-1}$ from R$_\mathrm{i}^{2}$ in \autoref{eq:mout} (where a factor of Z$_\mathrm{o}^{-2}$ comes directly from \zo in \autoref{eq:tau0}, and a factor of Z$^{+1}_\mathrm{o}$ comes from $\chi_\mathrm{SiIV}$ in \autoref{eq:zfrac}). Consequently, our measured \mout changes as $Z_\mathrm{o}^{-1}$, but the metal-loading factor remains {\it constant} because the metal-loading factor is the product \moutp$\times$\zop. For example, if the \zo of SBS~1415+437 hypothetically decreased from the modeled 0.96~Z$_\odot$ to the \zism\ of 0.08~Z$_\odot$, then \mout would increase from 0.3~\sfr to 3.6~\sfrp. This would increase the mass-loading factor from 19 to 225, but the {\it metal}-loading factor would remain 225. Therefore, if SBS~1415+437 did not drive an enriched outflow, then the outflow must remove {\it 12 times more} mass to match the observed MZR. Consequently, metal enriched outflows are {\it not} required to match the MZR, but if the outflows are not metal-enriched then the mass outflow rate must be sufficiently large enough to produce the required metal-loading factor. The outflow metallicity is an important new constraint for galaxy simulations, and may help to produce more realistic simulations. Typically, cosmological simulations drive galactic outflows with \zop~$ = $~\zism\ \citep{finlator08, dave11, schaye15, christensen, muratov17} or with $\mathrm{Z}_\mathrm{o} = 0.4 \mathrm{Z}_\mathrm{ISM}$ \citep{vogelsberger13}. The mass-loading factor corresponding to \zop~=~\zism\ for SBS1415+437 is 225, which is more consistent with the mass-loading factor of 140 predicted by the FIRE simulations \citep[eq. 4 in][]{muratov} than the mass-loading factor of 19 we estimated in \citetalias{chisholm17}. Thus, a simulation with \zop~=~\zism\ would drive 12~times more mass out of the galaxy than suggested by our observations. Simulations often have troubles producing starburst galaxies \citep{sparre15}, possibly because the simulations remove too much mass in order to match the observed MZRs. If galactic outflows are metal-enriched, then cosmological simulations that reproduce the observed MZR may drive too massive of galactic outflows and deplete their gaseous reservoir too rapidly. \subsubsection{How are outflows enriched?} \label{drive} \begin{figure} \includegraphics[width = 0.5\textwidth]{fe.pdf} \caption{The entrainment fraction ($f_e$; \autoref{eq:fe}) versus stellar mass (\mstarp). The entrainment fraction is the fraction of the outflow that is entrained ISM versus pure supernovae ejecta. An $f_e$ of 1.0 indicates that the outflow is entirely swept up ISM, while a lower $f_e$ indicates that the outflow is enriched by supernova ejecta. Outflows are largely swept up ISM, while outflows from low-mass galaxies have moderate enrichment.} \label{fig:fentrain} \end{figure} The outflow metallicity provides clues about the origins of the outflowing gas. A \zop/\zism~$=1$ would indicate that the outflow is entirely entrained ISM. Meanwhile, elevated \zop/\zism\ ratios imply that outflows are enriched by recent star formation (supernovae or stellar winds). Supernova ejecta have Z$_\mathrm{ej} \sim 0.1 = 7~Z_\odot$ \citep[][]{woosley95}, and this large supply of metals can mix with the ambient ISM to increase \zop. We observe that \zo is larger than \zism\ for six of the seven outflows (left panel of \autoref{fig:zout}), {and statistically consistent with being equal to \zism\ for one galaxy}, implying that the observed outflows are entrained ISM with some metal enrichment. We estimate the fraction of the outflow that is pure ISM. The "entrainment fraction" \citep[$f_\mathrm{e}$;][]{dalcanton, peeples11} is the fraction of the outflow that is purely swept up ISM, and is given as \begin{equation} f_e = \frac{Z_\mathrm{o}- Z_\mathrm{ej}}{Z_\mathrm{ISM} - Z_\mathrm{ej}} , \label{eq:fe} \end{equation} where we used $Z_\mathrm{ej} = 0.1$. An $f_\mathrm{e}$ of 1.0 indicates that the outflow is entirely swept up ISM, while an $f_e$ of zero indicates that the outflow is entirely supernova ejecta. Values between these extremes indicate that the ISM is a mixture of supernova ejecta and ambient ISM. High-mass galaxies have $f_\mathrm{e}$ values between 88 and 100 per cent (\autoref{fig:fentrain}); these outflows are nearly pure ISM. Even at the lowest stellar masses, 78 per cent of the outflowing gas from 1~Zw~18 is entrained ISM and only 22 per cent is pure supernova ejecta. Galactic outflows are largely entrained ISM with a small fraction of supernova enrichment. One possible origin of the $f_e$ behavior is that \zo of high-mass galaxies saturates as \zo approaches $y$. All outflows may have a flat \zop$\sim1$~Z$_\odot$, regardless of \zism\ or \mstar (see \autoref{fig:z}). A flat \zo causes $f_e$ to increase with increasing \zism\ until it reaches 1.0 at \zism~$=1$~Z$_\odot$ where it remains constant. A fixed \zo could arise if $y$ is fixed with \zism. This is compatible with current model yields which show that $y$ only changes by a factor of 1.06 from \zs\ of 0.05 to 2.0~Z$_\odot$, but varies strongly with different IMFs \citep[by a factor of $4-5$, see discussion above;][]{vincenzo}. Thus, an evolving IMF or an increasing $\zeta_0$ (removing more metals) are the most likely determinants of Z$_{\rm flat}$ (see \autoref{eq:zflat}), which leads to the lower Z$_{\rm flat}$ that are observed for higher redshift galaxies \citep{savaglio05, erb06}. Another possibility is that the observed outflows are enriched and driven differently in high and low-mass galaxies. These different physical driving mechanisms produce different \zop/\zism\ ratios. The smaller \zop/\zism\ of high-mass galaxies implies that these outflows largely remove ambient ISM material, which is expected if a supernova blast wave does not break out of the ISM, and instead "snowplows" ISM outward. Meanwhile, low-mass galaxies do not fully contain the supernovae blast wave, rather the ejecta rips holes in, and mixes with, the ambient ISM. Hints of this dichotomy are seen in the \siii and \siiv line profiles (\autoref{fig:s42}). \siii and \siiv are co-moving in high-mass galaxies, as expected in an outwardly expanding snowplow phase. Meanwhile, the \siii in lowest mass galaxies extends to redder velocities than the \siivp. This indicates that high-ionization gas in the outflows of low-mass galaxies extends to higher velocities than the low-ionization gas. Unlike the high-mass galaxies, the \siii of the two low-mass galaxies also has a slightly larger $C_f$ than the \siivp. If $C_f$ decreases with $r$ (\autoref{eq:cf}), this suggests that the \siii originates at smaller radii than the \siiv. This scenario is expected if low-ionization gas is heated as it accelerates, as expected if a supernova ripped a hole in the ambient ISM \citep{chisholm18}. \subsection{ {An examination of our assumptions}} {Here we discuss the implicit assumptions associated with our stellar continuum fitting, physical outflow model, photoionization model, chemical evolution model, and host galaxy property calculations have on determining the role galactic outflows play in shaping the MZR.} {In \autoref{cont} we modeled the observerd stellar continuum as a linear-combination of instantaneous bursts of fully-theoretical single-aged \textsc{starburst99} models. This modeling determines the number of ionizing photons and the hardness of the ionizing spectra, which we used in the \textsc{cloudy} photoionization modeling to determine the ionization structure and outflow metallicity. In Papers~\textsc{I}--\textsc{III} we demonstrated that the linear-combination of stellar continuum models reproduces the observed massive star features like the \ion{N}{V} and \siiv P-Cygni stellar wind features, however the assumed IMF and stellar atmospheric models impact the number of ionizing photons produced by a single burst \citep{claus99}. For instance, binary stellar populations have recently been suggested to have a harder ionizing spectra with more ionizing photons, which may reproduce observed features like \ion{He}{ii} emission \citep{eldridge, stanway}. The two galaxies in the sample with \ion{He}{ii}~1640\AA\ coverage (1~Zw~18 and NGC~7714) have mixed \ion{He}{ii}~1640\AA\ results. 1~Zw~18 has strong \ion{He}{ii}~1640\AA\ \citep{lebouteiller} and optical \ion{He}{ii}~4363\AA, although models show that the \ion{He}{ii}~4363\AA\ emission is even stronger than predicted by binary stars alone \citep{kerig}. Meanwhile, NGC~7714 does not have detected \ion{He}{ii}~1640\AA\ emission, illustrating that the hardness of the ionizing continua varies from galaxy-to-galaxy. Whether this is solely due to stellar metallicity affects, or requires a binary stellar population is still an unsolved problem.} {In \autoref{proffit} we fit for how the density and covering fraction scales with radius as well as the radial extent of the outflows. These parameters describe the acceleration and mass outflow rate. While we assume a physically motivated general form for the acceleration, density, and covering fraction; the exact scaling is observationally driven because we fit the parameters from the observed line profiles. This allows for galaxy-to-galaxy geometrical variations based on the observed outflow line profiles. For instance, the density profile sharply decreases with increasing radius, implying that galactic outflows are thin, but not arbitrarily thin, shells. The statistical uncertainties of each fitted parameter are included in the \mout uncertainties, which incorporates the fitted geometrical uncertainties into the \mz uncertainties. As discussed in \citetalias{chisholm17}, this is especially apparent in low-mass galaxies where narrow line profiles make it challenging to fit the line profile (see \autoref{tab:prof}).} {The photoionization modeling in \autoref{cloudy} assumes that the observed phases have similar physical conditions and are in photoionization equilibrium. Galaxies with stellar mass greater than 10$^{9}$~M$_\odot$ appear to satisfy these criteria {because the low and high-ionization lines have similar profiles} (\autoref{fig:s42}), but lower mass galaxies may not. As discussed in \autoref{drive}, this may indicate that there is unaccounted for physics in the production and acceleration of galactic outflows arising from low-mass galaxies, and, if confirmed by future observations, may lead to the generation of more realistic galactic outflows in low-mass galaxies.} {In \autoref{cloudy} we also assumed that the relative abundances do not vary from galaxy-to-galaxy. We used the \ion{H}{ii} abundance set in \textsc{CLOUDY} to set the relative abundance. Observations do not detect a trend in the relative abundances of $\alpha$-elements, like Si/O or S/O that we use here, with metallicity \citep{garnett95, izotov99}. This suggests that the assumption of constant relative abundance does not drastically impact the ionization models.} {One way to assess the impact of our assumed outflow model is to compare our \mout values to values found by other studies using similar data. \citet{heckman15} calculated \mout from NGC~7714 and Haro~11 assuming a constant outflow column density of 10$^{20.85}$~cm$^{-2}$, a constant outflow metallicity of 0.5~Z$_\odot$, and that all of the outflowing gas is at a constant radius that is twice the size of the star-forming region. The mass-loading factors differ by 0.02 and 0.46~dex between the two studies, within the errors of the measurement. The outflow assumptions do not heavily impact the derived mass-loading factors determined for these galaxies as compared to other studies. } {The fact that our empirical metal-loading factors agree with the analytic chemical evolution models of \citet{peeples11} supports the latter's conclusion that efficient outflows shape the MZR. However, one caveat is that these models assume that accreting gas has zero metallicity. Accretion is clearly an important process; 1~Zw~18 and Haro~11 lie below the MZR and their metallicities are often explained by recently accreted low-metallicity gas \citep{vanzee, ekta}. However, recent simulations show that most of the gas accreted onto galaxies at z~=~0 is not pristine gas, rather it is substantially metal-enriched, recycled galactic outflows \citep{angles}. If true, the metal-loading factors would have to be higher to match the observed MZR. More observational constraints on the accretion process are required to clarify the situation.} {We used the SFRs measured within the COS aperture as opposed to the global SFRs because we find that the bulk of the gas in the outflow is $< 100$~pc above the star-forming region (\autoref{tab:prof}), and thus more likely to be associated with local driving mechanisms rather than global ones. If we replaced the SFRs with global ones, the metal-loading factors would decrease by 0.27~dex on average. SFRs calculated from a combination of the UV and IR luminosities are more robust than SFRs calculated form the UV alone \citep{hao}. However, we do not have high-spatial resolution IR emission maps. Consequently, we assumed that the fraction of dust obscured star formation within the COS aperture is similar to the global average. However, since we preferentially target UV-bright regions which may have lower dust attenuation than the average, this may overestimate the SFR within the COS aperture, and thereby underestimate the mass and metal-loading factors. This would impact the two most massive galaxies in our sample the most (NGC 7714, NGC 6090) which have the highest dust attenuation (\autoref{tab:sample}).} {Both the UV and IR trace star formation on timescales of $\sim$100~Myr \citep{kennicutt2012}. Analysis of the resolved stellar populations in I~Zw~18 and SBS~1415+437 suggest that a large majority of the star formation has occurred over an extended period in the past 100~Myr \citep{mcquinn, annibali}, with up to a factor of 5 variation in amplitude during this time. This 100~Myr timescale is consistent with the timescale over which energy and momentum are injected into the gas by supernovae and stellar winds \citep{claus99}. Thus, our SFRs are appropriate to use in computing the mass and metal-loading factors of the outflow. We have used nebular oxygen abundance measurements from the star-forming regions closest to our COS aperture to define \zism. Thus our \zop/\zism\ values should not have any spatial or temporal biases.} \subsection{Future work and outstanding questions} \label{further} Above we presented results suggesting that the metal-loading factors of nearby galactic outflows are consistent with analytic requirements to shape the MZR (\autoref{fig:zout}). A large hurdle of this study is having a statistically relevant sample of star-forming galaxies covering the full range of galaxy properties. With only seven local galaxies, the small sample size cannot definitively discern whether galaxies have metal enriched galactic outflows or if galactic outflows shape the MZR. The small size is largely because it is observational challenging to observe the four weak ISM absorption lines due to contamination by Milky Way absorption and geocoronal emission lines at low redshifts. Crucially, the two low-mass galaxies in the sample have contamination from Milky Way absorption, geocoronal emission, and zero-velocity absorption (see \autoref{prev}). These low-mass galaxies also have the largest parameter uncertainties due to their weak and narrow line profiles (\autoref{fig:s42}). More high-quality observations of outflows from low-mass galaxies will determine the scaling of the metal-loading factor. This will provide new constraints for galaxy simulations by developing a new empirically motivated feedback constraint (see \autoref{highz}). While local galaxies are observable with COS, half of the total stars formed--and therefore metals produced--happened near $z \sim 2$. Therefore, these redshifts are crucial to understand why the shape of the MZR does not evolve with time, but the normalization does vary with redshift \citep{zahid}. Fortunately, at $z \sim 2$ the lines used in this analysis are redshifted into the optical, enabling a similar study of metal-loading factors in moderately redshifted galaxies. The advent of integral field spectroscopy (IFS) emphasizes that the metallicity depends both on the mass of the galaxy and the location within the galaxy. Metallicity gradients have been discovered in a variety of galaxy types in the local universe \citep{sanchez, belfiore}. Understanding how the outflow metallicity changes with spatial location within a galaxy will demonstrate the role outflows play in creating metallicity gradients. New, high through-put, IFS instruments, like MUSE and the Keck Cosmic Web Imager (KCWI), allow for similar measurements as we presented here to be made of spatially resolved gravitationally lensed galaxies at $z > 2$. Early spatially resolved galactic outflow studies are inconclusive whether outflows vary spatially \citep{bordoloi} or remain relatively constant within a galaxy \citep{james18}. Further results will determine whether spatially varying outflows produce metallicity gradients. Finally, the hot outflow phase is the least observed and most poorly constrained phase. The hot outflow phase is poorly constrained because current X-ray telescopes can only feasibly measure the outflow metallicities from a small sample of the nearest galaxies \citep{strickland2000, martin02, strickland09}. Future X-ray telescopes will have the collecting area and spectral resolution required to measure the metal outflows of a representative sample to determine the metallicity variation of the hot outflows. | 18 | 8 | 1808.10453 |
1808 | 1808.10179_arXiv.txt | Interstellar CO$_2$ is an important reservoir of carbon and oxygen, and one of the major constituents of the icy mantles of dust grains, but it is not observable directly in the cold gas because has no permanent dipole moment. Its protonated form, \HOCO , is believed to be a good proxy for gaseous CO$_2$. However, it has been detected in only a few star-forming regions so far, so that its interstellar chemistry is not well understood. We present new detections of \HOCO\ lines in 11 high-mass star-forming clumps. Our observations increase by more than three times the number of detections in star-forming regions so far. We have derived beam-averaged abundances relative to H$_2$ in between $0.3$ and $3.8\times 10^{-11}$. We have compared these values with the abundances of \HCOpI, a possible gas-phase precursor of \HOCO, and \METH, a product of surface chemistry. We have found a positive correlation with \HCOpI, while with \METH\ there is no correlation. We suggest that the gas-phase formation route starting from \HCOp\ plays an important role in the formation of \HOCO, perhaps more relevant than protonation of CO$_2$ (upon evaporation of this latter from icy dust mantles). | \label{intro} Carbon dioxide (CO$_2$) is a relevant molecular species in a variety of interstellar environments. In comets, planetary atmospheres, and interstellar ices, its abundance is a significant fraction ($\sim 0.1 - 0.5$) of that of water (e.g.~Bergin et al.~\citeyear{bergin}, Whittet et al.~\citeyear{whittet}, McKay et al.~\citeyear{mckay}, Hoang et al.~\citeyear{hoang}). CO$_2$ ice is one of the main constituent of the icy mantles of dust grains (\"Oberg et al.~\citeyear{oberg}). In the gas-phase, CO$_2$ can be observed directly through ro-vibrational transitions (e.g.~van Dishoeck et al.~\citeyear{vandishoeck}), but the lack of permanent dipole moment, and hence of a pure rotational spectrum, makes it impossible a detection in cold environments. Instead, its protonated form, \HOCO, has been detected towards the Galactic center (Thaddeus et al.~\citeyear{thaddeus}, Minh et al.~\citeyear{minh}, Neill et al.~\citeyear{neill}), in diffuse and translucent clouds (Turner et al.~\citeyear{turner}), but only in a handful of star-forming regions: in the low-mass pre-stellar core L1544 (Vastel et al.~\citeyear{vastel}), in the protostars L1527 and IRAS 16293--2422 (Sakai et al.~\citeyear{sakai}, Majumdar et al.~\citeyear{majumdar}), and in the protostellar shock L1157--B1 (Podio et al.~\citeyear{podio}). In cold and dense gas, two main chemical formation pathways have been proposed: {\bf (1)} a gas-phase route from the reaction ${\rm HCO^+ + OH \leftrightarrow HOCO^+ + H }$, and {\bf (2)} the protonation of CO$_2$ (mainly upon reaction with H$_3^+$) desorbed from grain mantles (see e.g.~Vastel et al.~\citeyear{vastel}, Bizzocchi et al.~\citeyear{bizzocchi}). In scenario {\bf (1)}, CO$_2$ would be a product of \HOCO\ (after dissociative recombination), while the opposite is expected in scenario {\bf (2)}. Due to the lack of stringent observational constraints, it is unclear yet which of these two mechanisms is dominant, and under which physical conditions. Constraining the abundance of \HOCO\ has important implications also for the abundance of CO$_2$ in ice. In fact, if \HOCO\ is formed in the cold gas and then, upon dissociative recombination, gives rise to CO$_2$, this latter could freeze-out on grain mantles and contribute to the amount of CO$_2$ ice observed in dark clouds (Bergin et al.~\citeyear{bergin}), although this cannot explain the large amount of solid CO$_2$ measured along the line of sight of background stars (Boogert et al.~\citeyear{boogert}) or deeply embedded massive young stars (van Dishoeck et al.~\citeyear{vandishoeck}). In fact, the formation of CO$_2$ ice from surface reactions is still debated. Laboratory experiments suggested formation of CO$_2$ ice from ${\rm CO + O \rightarrow CO_2}$ (D'Hendecourt et al.~\citeyear{dhendecourt}), which however needs a strong UV irradiation, and hence it is expected to be inefficient in dark clouds. Other surface reactions have been proposed, such as cosmic-ray bombardment on carbonaceous grains covered by water ice (Mennella et al.~\citeyear{mennella}), or the radical-radical reaction ${\rm OH + CO \rightarrow CO_2 + H}$ (Garrod \& Pauly~\citeyear{gep}, Ioppolo et al. 2011, Noble et al. 2011). However, such process involves the diffusion of heavy radicals, difficult to happen at dust temperatures below $\sim 30$~K, although high precision measurements are not available yet, and new promising techniques have been proposed (e.g.~Cooke et al.~\citeyear{cooke}) to shed light on this important surface chemistry process. In this paper, we present new detections of \HOCO\ in 11 high-mass star-forming regions, belonging to an evolutionary sample of 27 clumps divided into the three main evolutionary categories of the massive star formation process (Fontani et al.~\citeyear{fontani2011}): high-mass starless cores (HMSCs), high-mass protostellar objects (HMPOs) and Ultra-compact \HII\ regions (UCHIIs). The sample has been extensively observed in several dense gas tracers, with the aim of studying the chemical evolution of these molecules during the massive star-formation process (Fontani et al.~\citeyear{fontani2011},~\citeyear{fontani2014},~\citeyear{fontani2015a},~\citeyear{fontani2015b}, ~\citeyear{fontani2016}, Colzi et al.~\citeyear{colzi},~\citeyear{colzib}, Mininni et al.~\citeyear{mininni}). This work represents the first study of protonated carbon dioxide in a statistically relevant number of star forming regions. | \label{discu} \begin{figure} \begin{center} \includegraphics[width=6.3cm,angle=0]{abundances_plots_n2hp_scuba_new.eps} \caption{Abundance of \HOCO\ against that of, from top to bottom: \METH, H$^{13}$CO$^+$, and \H. The colours indicate the different evolutionary groups as labelled in the top-left corner. The large filled symbols correspond to the detected sources, while the small triangles indicate the upper limits on the abundance of \HOCO. In the panel with X[H$^{13}$CO$^+$], we do not show 05358-mm3, observed and detected in \METH, \H\ and \HOCO\ but not in H$^{13}$CO$^+$.} \label{fig_abb} \end{center} \end{figure} \begin{figure} \begin{center} {\includegraphics[width=8.1cm,angle=0]{corr-linew-ch3oh-hocop.eps}} \caption{Line widths at half maximum, $\Delta v$, of \HCOpI\ (left panel) and \METH\ (right panel) against those of \HOCO. The colours indicate the different evolutionary groups as in Fig.~\ref{fig_abb}. The solid line in the left panel corresponds to a linear fit to the data.} \label{fig_dv} \end{center} \end{figure} As discussed in Sect.~\ref{intro}, two main pathways have been proposed for the formation of \HOCO\ in dense gas: either the gas-phase reaction ${\rm HCO^+ + OH \leftrightarrow HOCO^+ + H }$, or the protonation of desorbed CO$_2$, mainly upon reaction with H$_3^+$ (Vastel et al.~\citeyear{vastel}, Bizzocchi et al.~\citeyear{bizzocchi}). To investigate if and how our observational results can put constraints on these alternative pathways, in Fig.~\ref{fig_abb} we show the fractional abundance of \HOCO\ (calculated as explained in Sect.~\ref{abundances}) against that of \HCOpI, the main precusor of \HOCO\ in the gas phase, and \METH, a tracer of surface chemistry. For completeness, we also plot X[HOCO$^+$] against X[N$_2$H$^+$], because protonation of CO$_2$ may occur also via reaction with \H. The \METH\ and \H\ column densities used to compute X[CH$_3$OH] and X[N$_2$H$^+$] have been taken from Fontani et al.~(\citeyear{fontani2015a}) and Fontani et al.~(\citeyear{fontani2015b}), respectively, and rescaled to the beam of the \HOCO\ observations. The \HCOpI\ column densities used to derive X[H$^{13}$CO$^+$] have been estimated from the integrated intensities of the \HCOpI\ 1--0 lines at 86754.288~MHz, serendipitously detected in the same dataset described in Sect.~\ref{obs} and in Colzi et al.~(\citeyear{colzi}). We have followed the same approach used for \HOCO, namely we assumed optically thin lines and LTE conditions (see Sect.~\ref{abundances}). We used the excitation temperatures listed in Table~1. The beam size is almost the same of that of \HOCO, hence all the fractional abundances in Table~1 are averaged over the same angular region. Figure~\ref{fig_abb} indicates a clear non-correlation between the abundances of \METH\ and \HOCO, while \HCOpI\ and \HOCO\ seems positively correlated. By applying simple statistical tests to the detected sources only, the correlation coefficient (Pearson's $\rho$) between X[\HOCO] and X[\HCOpI] is 0.7. Considering also the upper limits, the correlation remains positive (Pearson's $\rho$ = 0.6). The correlation between X[\HOCO] and X[\H] is positive (Pearson's $\rho$ = 0.4 without the upper limits) but much less convincing. If we assume that both CO$_2$ and \METH\ form on grain mantles, and what we find in the gas is evaporated at similar times, the lack of correlation between X[\HOCO] and both X[\METH] and X[\H] would indicate that the origin of \HOCO\ is likely not from CO$_2$ evaporated from ice mantles. This interpretation has two big caveats. First, the formation processes of CO$_2$ and \METH\ on the surfaces of dust grains can be different. In fact, CO$_2$ is thought to form in water ice mantles of cold carbonaceous grains via cosmic-ray bombardment (Mennella et al.~\citeyear{mennella}), or via the surface reaction CO + OH at dust temperatures of $\sim $30~K (Garrod \& Pauly~\citeyear{gep}), while \METH\ if formed from hydrogenation of CO at dust temperatures $\sim 10$~K (Vasyunin et al.~2017). Second, the main molecular ion responsible for the protonation of CO$_2$ is H$_3^+$ and not \H. Nevertheless, the positive correlation between X[\HCOpI] and X[\HOCO] suggests a non negligible, or even dominant, contribution from \HCOp\ to the formation of the detected \HOCO. This interpretation of our results is in agreement with the study of Majumdar et al.~(\citeyear{majumdar}), who proposed that the dominant (up to 85$\%$) formation route of \HOCO\ in the extended and cold ($T\leq 30$~K) envelope of the hot corino IRAS 16293--2422 is indeed from the gas-phase reaction OH + \HCOp. The fact that \HOCO\ and \HCOpI\ likely arise from similar gas can be understood also from the comparison of their line widths at half maximum. Fig.~\ref{fig_dv} indicates that the \HOCO\ line widths are correlated with those of \HCOpI, but not with those of \METH\ (which are always narrower). Hence, \METH\ is likely not associated with the same gas. Overall, our findings suggest a significant (perhaps dominant) role of HCO$^+$ as a gas-phase progenitor of \HOCO. However, caution needs to be taken in the interpretation of our results for several reasons. First, the \HOCO\ column densities, and hence the fractional abundances, have been derived assuming an excitation temperature that could not be that of the molecule. To solve this problem, detection of more lines tracing different excitation conditions are absolutely required. Another caveat arises from the fact that our column densities are values averaged over large (28\asec) angular surfaces. Our targets are known to have complex structure, and temperature (and density) gradients. Therefore, higher angular resolution observations are needed to precisely determine the \HOCO\ emitting region, and, from this, understand its temperatures and densities, required to properly model the chemistry. {\it Acknowledgments.} We thank the IRAM-30m staff for the precious help during the observations. We thank the anonymous referee for his/her constructive comments. V.M.R. and M.P. acknowledge the financial support received from the European Union's Horizon 2020 research and innovation programme under the Marie Sklodowska-Curie grant agreement No 664931. L.C. acknowledges support from the Italian Ministero dell'Istruzione, Universit\`a e Ricerca through the grant Progetti Premiali 2012 - iALMA (CUP C52I13000140001). P.C. acknowledges support from the European Research Council (ERC project PALs 320620). \let\oldbibliography\thebibliography \renewcommand{\thebibliography}[1]{\oldbibliography{#1} \setlength{\itemsep}{-1pt}} | 18 | 8 | 1808.10179 |
1808 | 1808.02136_arXiv.txt | We study the total and dark matter (DM) density profiles as well as their correlations for a sample of 15 high-mass galaxy clusters by extending our previous work on several clusters from Newman et al. Our analysis focuses on 15 CLASH X-ray-selected clusters that have high-quality weak- and strong-lensing measurements from combined Subaru and {\em Hubble Space Telescope} observations. The total density profiles derived from lensing are interpreted based on the two-phase scenario of cluster formation. In this context, the brightest cluster galaxy (BCG) forms in the first dissipative phase, followed by a dissipationless phase where baryonic physics flattens the inner DM distribution. This results in the formation of clusters with modified DM distribution and several correlations between characteristic quantities of the clusters. We find that the central DM density profiles of the clusters are strongly influenced by baryonic physics as found in our earlier work. The inner slope of the DM density for the CLASH clusters is found to be flatter than the Navarro--Frenk--White profile, ranging from $\alpha=0.30$ to $0.79$. We examine correlations of the DM density slope $\alpha$ with the effective radius $R_\mathrm{e}$ and stellar mass $M_\mathrm{e}$ of the BCG, finding that these quantities are anti-correlated with a Spearman correlation coefficient of $\sim -0.6$. We also study the correlation between $R_\mathrm{e}$ and the cluster halo mass $M_{500}$, and the correlation between the total masses inside 5\,kpc and 100\,kpc. We find that these quantities are correlated with Spearman coefficients of $0.68$ and $0.64$, respectively. These observed correlations are in support of the physical picture proposed by Newman et al. | The $\Lambda$ cold dark matter ($\Lambda$CDM) paradigm gives a plethora of correct predictions \citep{Komatsu2011,Planck2014,DelPopolo2013,DelPopolo2007,DelPopolo2014aa}. However, some of its predictions are at odds with observations. $N$-body simulations in $\Lambda$CDM predict that the spherically averaged density profiles of self-gravitating structures, ranging from dwarf galaxies to galaxy clusters, are cuspy and well approximated by the Navarro--Frenk--White (NFW) profile \citep{Navarro1997,Navarro2010}. However, observations \citep{Moore1994,Flores1994,Agnello2012,Adams2014} and theoretical studies \citep{Navarro1996,DelPopolo2009,Governato2010,DelPopolo2010,DelPopolo2011,DelPopolo2012a} have shown that the inner slopes of the density profile in dwarf galaxies and low-surface-brightness galaxies (LSBs) are usually flatter than simulations, and there is a strong diversity of the dark-matter (DM) distribution in these low-mass systems \citep[the so-called ``diversity problem'',][]{Simon2005,DelPopolo2012a,Oman2015}.\footnote{In addition to this problem, the $\Lambda$CDM paradigm suffers from the cosmological constant problem \citep{Weinberg1989,Astashenok2012}, the unknown nature of dark energy \citep{DelPopolo2013a,DelPopolo2013b,DelPopolo2013c}, and from several problems at small scales \citep{DelPopolo2017,DelPopolo2017a}} On the observational side, the small dynamic range of observations can cause a degeneracy in the mass profile determination \citep[see][]{DelPopolo2002}, and this degeneracy cannot be fully broken due to the lack of HI observations in dwarf spheroidals (dSPhs) and elliptical galaxies. Determinations of their DM structure are thus much more complicated. In the case of dSPhs, there are discrepant results on the cusp-core nature of the density profile \citep{Amorisco2012,Jardel2012,Jardel2013a,Jardel2013b}, sometimes even in the case of the same object studied with different techniques. Similar uncertainties are present in cluster of galaxies, but X-ray observations, lensing and galaxies dynamics overcome them in an easier manner than for the cases of dSPhs or ellipticals. While the NFW profile \citep{Navarro1996,Navarro1997} describes well the observed total density profiles in galaxy clusters as found in several studies \citep{Sand2002,Sand2004,Newman2009,Umetsu2011,Newman2011,Okabe2013,Newman2013a,Newman2013b,Umetsu2014,Umetsu2016}, it was also found that the inner DM structure is characterized by a flatter slope within typical scales of the brightest cluster galaxy (BCG; from some kpcs to some tens of kpcs). Hence, the cusp-core problem \citep{DelPopolo2012a,Oman2015} appears to be present in galaxy clusters as well. This discrepancy can be alleviated when the effects of baryonic physics are properly accounted for in $N$-body simulations \citep[see][]{ElZant2004,Nipoti2004,DelPopolo2009,DelPopolo2009a,DelPopolo2010,Cardone2011a,Cardone2011b,Governato2010,Cole2011, DelPopolo2011,Martizzi2012,Martizzi2013,Nipoti2015}. \begin{table}[htbp] \centering \caption{Parameters derived for the CLASH sample. First column: cluster name; second: $M_{500}$ as given in \citep{Umetsu2016}; third and fourth: innermost 2D density slopes inferred directly from the observed \citep{Umetsu2016} profiles and obtained from our semi-analytical model; fifth: inner 3D density slope from our model; sixth and seventh: stellar and baryonic fractions from our model. } \small \begin{tabular}{lllllll} \hline Name & $M_{500}$ & $\alpha_\mathrm{2D}$ & $\alpha_\mathrm{2D,T}$ & $\alpha_\mathrm{3D}$ & $f_\mathrm{star}$ & $F_\mathrm{b}$ \\ & [$10^{14}M_\odot$] & & & & & \\ \hline A383 & 5.88 $\pm$1.73 & 0.71 $\pm$0.26 & 0.70 $\pm$0.09 & 0.37 $\pm$0.09 & 0.0201 $\pm$0.002 & 0.1355 $\pm$0.008 \\ A209 & 9.64 $\pm$1.97 & 0.67 $\pm$0.29 & 0.68 $\pm$0.09 & 0.60 $\pm$0.1 & 0.0167 $\pm$0.002 & 0.1417 $\pm$0.01 \\ A2261 & 15.65 $\pm$3.05 & 0.77 $\pm$0.26 & 0.79 $\pm$0.09 & 0.63 $\pm$0.09 & 0.0140 $\pm$0.002 & 0.1480 $\pm$0.012 \\ RXJ2129 & 4.48 $\pm$1.16 & 0.49 $\pm$0.26 & 0.49 $\pm$0.09 & 0.55 $\pm$0.09 & 0.0222 $\pm$0.002 & 0.1323 $\pm$0.006 \\ A611 & 10.73 $\pm$2.65 & 0.59 $\pm$0.27 & 0.58 $\pm$0.09 & 0.79 $\pm$0.09 & 0.0161 $\pm$0.002 & 0.1431 $\pm$0.01 \\ MS2137 & 8.28 $\pm$2.57 & 0.86 $\pm$0.25 & 0.85 $\pm$0.09 & 0.65 $\pm$0.08 & 0.0177 $\pm$0.002 & 0.1398 $\pm$0.009 \\ RXJ2248 & 12.45 $\pm$3.62 & 0.45 $\pm$0.28 & 0.44 $\pm$0.09 & 0.55 $\pm$0.09 & 0.0152 $\pm$0.002 & 0.1450 $\pm$0.011 \\ MACSJ1115 & 10.67 $\pm$2.22 & 0.33 $\pm$0.30 & 0.34 $\pm$0.09 & 0.39 $\pm$0.1 & 0.0161 $\pm$0.002 & 0.1430 $\pm$0.01 \\ MACSJ1931 & 10.51 $\pm$4.05 & 0.69 $\pm$0.28 & 0.70 $\pm$0.09 & 0.65 $\pm$0.09 & 0.0162 $\pm$0.002 & 0.1428 $\pm$0.01 \\ MACSJ1720 & 9.96 $\pm$2.53 & 0.59 $\pm$0.26 & 0.60 $\pm$0.09 & 0.56 $\pm$0.09 & 0.0165 $\pm$0.002 & 0.1421 $\pm$0.01 \\ MACSJ0429 & 6.85 $\pm$2.1 & 0.45 $\pm$0.28 & 0.44 $\pm$0.09 & 0.48 $\pm$0.09 & 0.0190 $\pm$0.002 & 0.1374 $\pm$0.008 \\ MACSJ1206 & 12.24 $\pm$2.49 & 0.56 $\pm$0.26 & 0.57 $\pm$0.09 & 0.50 $\pm$0.09 & 0.0153 $\pm$0.002 & 0.1448 $\pm$0.011 \\ MACSJ0329 & 6.51 $\pm$1.37 & 0.65 $\pm$0.27 & 0.64 $\pm$0.09 & 0.70 $\pm$0.09 & 0.0193 $\pm$0.002 & 0.1368 $\pm$0.008 \\ RXJ1347 & 22.33 $\pm$4.89 & 0.39 $\pm$0.30 & 0.40 $\pm$0.09 & 0.30 $\pm$0.1 & 0.0123 $\pm$0.002 & 0.1528 $\pm$0.014 \\ MACSJ0744 & 11.94 $\pm$2.81 & 0.54 $\pm$0.27 & 0.53 $\pm$0.09 & 0.55 $\pm$0.09 & 0.0155 $\pm$0.002 & 0.1445 $\pm$0.011 \\ \hline \end{tabular}% \label{tab:1}% \end{table}% \begin{table}[htbp] \centering \caption{Physical parameters derived for the CLASH sample. First column: cluster name; second: BCG mass derived from our model; third: BCG effective radius; fourth and fifth: spherical total masses inside 5\,kpc and 100\,kpc.} \small \begin{tabular}{lllll} \hline Name & $M_\mathrm{e}$ & $R_\mathrm{e}$ & $M_\mathrm{5 kpc}$ & $M_\mathrm{100 kpc}$\\ & [$10^{11}M_\odot$] & [kpc] & [$10^{11}M_\odot$] & [$10^{13}M_\odot$]\\ \hline A383 & 9.16 $\pm$0.29 & 28.7 $\pm$1.5 & 0.98 $\pm$0.15 & 1.96 $\pm$0.3 \\ A209 & 7.84 $\pm$0.29 & 25 $\pm$1.5 & 1.31 $\pm$0.15 & 3.21 $\pm$0.3 \\ A2261 & 10.5 $\pm$0.29 & 40 $\pm$1.5 & 1.75 $\pm$0.15 & 4.32 $\pm$0.3 \\ RXJ2129 & 13.73 $\pm$0.29 & 33 $\pm$1.5 & 2.43 $\pm$0.15 & 3.49 $\pm$0.3 \\ A611 & 12.25 $\pm$0.29 & 34.6 $\pm$1.5 & 1.63 $\pm$0.15 & 3 $\pm$0.3 \\ MS2137 & 6.55 $\pm$0.29 & 14 $\pm$1.5 & 1.95 $\pm$0.15 & 3.5 $\pm$0.3 \\ RXJ2248 & 12.45 $\pm$0.29 & 38.5 $\pm$1.5 & 2.17 $\pm$0.15 & 4.3 $\pm$0.3 \\ MACSJ1115 & 11.44 $\pm$0.29 & 44.5 $\pm$1.5 & 2 $\pm$0.15 & 3.56 $\pm$0.3 \\ MACSJ1931 & 8.19 $\pm$0.29 & 31 $\pm$1.5 & 1.37 $\pm$0.15 & 3.5 $\pm$0.3 \\ MACSJ1720 & 8.99 $\pm$0.29 & 35.8 $\pm$1.5 & 1.5 $\pm$0.15 & 3.32 $\pm$0.3 \\ MACSJ0429 & 13.37 $\pm$0.29 & 41 $\pm$1.5 & 2.32 $\pm$0.15 & 3.28 $\pm$0.3 \\ MACSJ1206 & 11.96 $\pm$0.29 & 43 $\pm$1.5 & 2.08 $\pm$0.15 & 4.08 $\pm$0.3 \\ MACSJ0329 & 7.41 $\pm$0.29 & 20 $\pm$1.5 & 1.24 $\pm$0.15 & 2.17 $\pm$0.3 \\ RXJ1347 & 13.8 $\pm$0.29 & 46.9 $\pm$1.5 & 2.45 $\pm$0.15 & 4.5 $\pm$0.3 \\ MACSJ0744 & 10 $\pm$0.29 & 37.1 $\pm$1.5 & 1.67 $\pm$0.15 & 3.98 $\pm$0.3 \\ \hline \end{tabular}% \label{tab:2}% \end{table}% In order to study how baryons modify the formation and evolution of clusters, we consider in \citep{DelPopolo2012b} baryonic clumps interacting with the DM model introduced in \citep{DelPopolo2009}. In addition to finding that the central baryonic concentration within 10\,kpc plays an important role in shaping the cluster density profile, we reproduced the observed cluster profiles for several massive systems \citep{Sand2004,Newman2009,Newman2011}, namely A611, A383, MACSJ1423.8+2404, and RXJ1133. In \citep{DelPopolo2014}, we reproduced the correlations found by \citep{Newman2013a,Newman2013b},\footnote{The quoted authors found correlations of the inner slope of the DM profile with the size of the BCG, the core radius, namely the constant density core of the cored NFW density profile \citep[see Eq.~2 of][]{Newman2013b}, and the BCG mass, and finally the correlation between the masses contained inside 5\,kpc and 100\,kpc. {% In the present paper, with BCG mass we refer to the stellar mass only.}} for MS2137, A963, A383, A611, A2537, A2667, and A2390. For these clusters, the total mass density profiles are well fitted by an NFW profile, while the central DM distribution is shallower than the total mass distribution. The formation picture proposed by \citep{Newman2013a,Newman2013b} is characterized by a dissipational formation of BCGs, followed by a dissipationless phase. In this phase, as described by \citep{ElZant2004,Ma2004,Nipoti2004,RomanoDiaz2008,RomanoDiaz2009,DelPopolo2009,Inoue2011, DelPopolo2012a,DelPopolo2012b,Cole2011,Nipoti2015}, baryon clumps interact with DM through dynamical friction, ``heating'' DM and reducing the central cusp. Our aims here are to use high-quality gravitational lensing observations from the CLASH survey \citep{Postman2012} and investigate if CLASH clusters exhibit correlations that are similar to those observed in the \citep{Newman2013a,Newman2013b} clusters, to characterize the mass distributions of CLASH clusters, and to test the physical picture that was proposed by \citep{Newman2013a,Newman2013b} and confirmed by \citep{DelPopolo2014}. To this end, we perform an improved analysis on a sample of 15 X-ray-selected CLASH clusters compared to our previous work \citep{DelPopolo2012b,DelPopolo2014}. In the present work, we will characterize the total mass density profiles of 15 CLASH clusters by means of a modified version of the semi-analytical model developed by \citep{DelPopolo2012b,DelPopolo2014}. Here we take into account the following effects: \begin{enumerate} \item adiabatic contraction (AC) responsible for the steepening of the inner density profiles in the early stage of cluster formation, \item interaction between baryonic clumps and DM through dynamical friction, which is responsible for ``heating'' the DM component and flattening the density profile, \item supernovae (SN) feedback, \item AGN feedback and other baryonic effects described in detail in Appendix. \end{enumerate} The paper is organized as follows. In section \ref{sec:DataModel}, we describe the data used and provide a brief summary of our model. In section \ref{sec:Results}, we discuss the results, and section \ref{sec:Conclusions} is devoted to conclusions. Throughout this paper, we adopt a concordance $\Lambda$CDM cosmology with $\Omega_\mathrm{m}=0.27$, $\Omega_\Lambda=0.73$, and $h=0.7$ with $H_0=100h$\,km\,s$^{-1}$\,Mpc$^{-1}$. | \label{sec:Conclusions} In this paper, we have studied and characterized the total density profiles for a sample of 15 X-ray-selected CLASH clusters by improving our earlier analysis \citep{DelPopolo2014} based on several clusters from \citep{Newman2013a,Newman2013b}. The primary goal of this study was to test the physical picture of cluster formation proposed by \citep{Newman2013a,Newman2013b} in the frame work of a modified version of the physical model developed by \citep{DelPopolo2012b,DelPopolo2014}. To this end, we analyzed binned surface mass density profiles of \citep{Umetsu2016} derived from their strong-lensing, weak-lensing shear and magnification analysis of high-quality {\em HST} and Subaru data. For each cluster, we extracted the radial profile of the total 3D density assuming spherical symmetry. We have used our semi-analytical model to interpret the total 3D density profile, which allows us to compute the baryon density profile and thus the DM density profile for the cluster. The total 3D mass density profile for our sample is characterized by a logarithmic slope of $\langle\gamma_\mathrm{tot}\rangle =1.05 \pm 0.02$ in the radial range $r=(0.003-0.03) \times r_{200}$, in agreement with several previous studies (e.g. \citep{Newman2013a,Newman2013b}). Stellar mass dominates the total mass at $r\lesssim 5-10$\,kpc, while the cluster outskirts are dominated by DM. Such segregation reveals a ``tight coordination'' between the inner DM and stellar distributions, as also implied by interplay of DM and baryons that generates the NFW-like total density profile. The correlation between the mass inside 5\,kpc and that inside 100\,kpc (Fig.~\ref{fig:6}) further supports such a tight coordination and points to similar formation time-scales of the BCG and the inner cluster region. Thus, the cluster's final configuration depends on the baryonic content and their formation process \citep{DelPopolo2012a}. Therefore, in the context of hierarchical structure formation models, we should expect tight correlations between the final inner baryonic content and the BCG mass, as well as between the total baryonic and cluster masses \citep[see][]{Whiley2008}. Since the DM and baryon contents sum to the total mass and the baryons dominate the inner $5-10$\,kpc region, the inner slope of the DM density must be shallower than that of the total density, that is, $\alpha<1$, as shown in Fig.~\ref{fig:1}. The observed inner DM slopes $\alpha$ span the range $[0.30, 0.79]$. Correlations were also examined between several of the characteristic quantities of clusters (e.g., $R_\mathrm{e}$, $M_\mathrm{e}$, $M_{500}$). Our findings are summarized as follows: \renewcommand{\theenumi}{\alph{enumi}} \begin{enumerate} \item The inner 3D slope of the DM density, $\alpha$, is anti-correlated with the BCG effective radius $R_\mathrm{e}$. The anti-correlation reflects the balance between DM and the BCG in the cluster center. For an NFW-like total mass profile, clusters with more massive BCGs contain less central DM, implying a flatter DM slope. \item Similarly, the inner DM slope $\alpha$ and the BCG mass $M_\mathrm{e}$ are anti-correlated with each other. This indicates again that a larger content of the central baryons gives rise to flatter DM profiles. \item The cluster halo mass $M_{500}$ and the BCG effective radius $R_\mathrm{e}$ are correlated with each other, as found in previous studies \citep{Kravtsov2013}. \item The cluster mass inside 5\,kpc, dominated by the stellar baryons, and the cluster mass inside 100\,kpc, dominated by DM, are correlated with each other. This hints at early formation of the BCG and the inner cluster region, while subsequent, continuous mass accretion played a fundamental role in the growth of cluster outskirts \citep[e.g.,][]{Fujita2018a,Fujita2018b}. \end{enumerate} These observed correlations are in support of the physical picture proposed by \citep{Newman2013a,Newman2013b}, that clusters form from a dissipative phase that leads to steepening the central stellar density, followed by a second dissipationless phase in which interactions between baryonic clumps and DM through dynamical friction kinematically heat the latter, leading to flat DM density profiles \citep{ElZant2001,ElZant2004,RomanoDiaz2008,DelPopolo2012a,DelPopolo2012b,Cole2011,Nipoti2015}. | 18 | 8 | 1808.02136 |
1808 | 1808.09823_arXiv.txt | {Two exotic elements have been introduced into the standard cosmological model: non-baryonic dark matter and dark energy. The success in converting a hypothesis into a solid theory depends strongly on whether we are able to solve the problems in explaining observations with these dark elements and whether the solutions of these problems are unique within the standard paradigm without recourse to alternative scenarios. We have not achieved that success yet because of numerous inconsistencies, mainly on galactic scales, the non-detection so far of candidate particles for dark matter, and the existence of many alternative hypotheses that might substitute the standard picture to explain the cosmological observations. A review of some ideas and facts is given here.} | 18 | 8 | 1808.09823 |
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1808 | 1808.06629_arXiv.txt | We present a library of high-resolution (R $\equiv$ \lam/$\Delta$\lam ~$\sim$ 45,000) and high signal-to-noise ratio (S/N $\geq$~200) \NIR spectra for stars of a wide range of spectral types and luminosity classes. The spectra were obtained with the Immersion GRating INfrared Spectrograph (IGRINS) covering the full range of the H (1.496-1.780~\um) and K (2.080-2.460~\um) atmospheric windows. The targets were primarily selected for being MK standard stars covering a wide range of effective temperatures and surface gravities with metallicities close to the Solar value. Currently, the library includes flux-calibrated and telluric-absorption-corrected spectra of 84 stars, with prospects for expansion to provide denser coverage of the parametric space. Throughout the H and K atmospheric windows, we identified spectral lines that are sensitive to \teff\ or \logg\ and defined corresponding spectral indices. We also provide their equivalent widths. For those indices, we derive empirical relations between the measured equivalent widths and the stellar atmospheric parameters. Therefore, the derived empirical equations can be used to calculate \teff\ and \logg\ of a star without requiring stellar atmospheric models. | \label{sec:intro} Spectral libraries are important for science efficiency and repeatability. By understanding standard stars very well, we can also study stellar populations and chemical abundances throughout the Galaxy. While a number of stellar spectral libraries are available, many are at optical wavelengths and low- to moderate-resolution. The number of spectral libraries with high spectral resolution in the \NIR (NIR) is very limited. NIR spectral libraries are useful for studying the cool phenomena of the universe: physics of cool stars, circumstellar objects such as planets and disks surrounding young stellar objects (YSOs), and the extended atmospheres of evolved stars. Compared to the red part of their optical spectra, the NIR spectra of cool stars are better suited for detailed spectroscopic study (especially in the J and H atmospheric windows, hereafter bands) as they contain fewer molecular bands. When determining equivalent widths (EWs) the broad molecular bands make the continuum level hard to determine and affect the measurement of narrow atomic lines. Also, the smaller extinction in the NIR, relative to optical wavelengths, provides us access to dust obscured stars across the Galaxy. Additionally, a NIR stellar library is useful for population synthesis of cool stars at the wavelengths where they are brightest. Over the past four decades, a number of infrared spectral libraries have been developed (see Table 1 in Rayner et al. 2009). For example, \citet{rayner09} presented an intermediate-resolution (R $\equiv $ \lam/$\Delta$\lam ~$\sim$ 2000) NIR spectral library of 210 F, G, K, and M stars that covers 0.8-5~\um, and this library provided flux calibrated spectra across this wide wavelength range for the first time. Moderate-resolution spectra are suitable for studying variations in strong spectral features, but many lines remain blended. For example, Figure~\ref{comp_spec} shows how the important Na I features at 2.2~\um\ are contaminated by lines of Si and Sc when observed at lower resolution \citep[see Table 7 in][]{ramirez97, doppmann03}. High-resolution spectroscopy \textquotedblleft resolves\textquotedblright~this problem. In addition, \citet{lebzelter12} provided a library of high-resolution (R $\sim$ 100,000) NIR spectra for 25 stars observed with CRIRES at the VLT. However, the small number of stars observed with CRIRES leaves the spectral type and luminosity class space coarsely sampled. The primary reason for the small sample in the CRIRES library is that the broad spectral coverage (1-5~\um) required about 200 instrument settings to cover the full wavelength range, and about 70 settings to cover the entire H- and K-bands. Other large surveys like APOGEE \citep[][and references therein]{zamora15} will obtain hundreds of thousands of NIR spectra, but the planned spectral coverage is narrow and it will not provide a comprehensive library covering the entire parametric space. Currently, a uniform library of high-resolution NIR spectra covering all spectral types and luminosity classes is not available. Although low- and moderate-resolution spectra are appropriate for extragalactic studies, because of the faintness of the targets and their internal velocity dispersions of 100 to 250~\kms, there is an ongoing need for a high-resolution spectral libraries in the Galactic context. One consequence of the absence of such libraries is that only strong and isolated lines have been used for studies. High-resolution NIR spectra will reveal lines with small equivalent widths that can be more sensitive to stellar properties and to abundances of less common elements \citep{afsar16}. At R=45,000, lines in most stars are well-resolved and the line shapes offer an additional window into the stellar properties. High-resolution stellar spectra covering wide ranges of effective temperature and surface gravity will test atmospheric models with unprecedented accuracy and permit modeling to improve infrared atomic and molecular line data. In this paper, we present a high-resolution (R $\sim$ 45,000) and high signal-to-noise ratio (S/N $\geq$~200) NIR spectral library covering a broad range of spectral parameters. The spectra were obtained with the Immersion GRating INfrared Spectrograph \citep[IGRINS,][]{yuk10, park14, mace16}. IGRINS has a resolving power of R $\sim$ 45,000 and covers the full H (1.49-1.80~\um) and K (1.96-2.46~\um) bands, simultaneously. The data presented in this paper was obtained as part of the \textquotedblleft IGRINS spectral library (PI : Jeong-Eun Lee)\textquotedblright~Legacy program\footnote{Korean IGRINS Legacy Program - \\ \url{http://kgmt.kasi.re.kr/kgmtscience/kgmt_igrins/legacy/selected.html}}. We present spectra, derived equivalent widths, and empirical relations based on the spectral catalog. | \label{sec:conclusions} We present the IGRINS Spectral Library of NIR spectra for 84 stars with high-resolution (R $\sim$ 45,000) and high signal-to-noise ratio (S/N $\geq$~200), as part of the IGRINS Legacy Project. The library covers O- to M-type stars with luminosity classes between I and V. The observed spectra were reduced using the IGRINS pipeline version 2, and hydrogen lines were removed by fitting the Kurucz Vega model. Then, we performed a correction for the telluric absorption lines and absolute flux calibration using the 2MASS photometry and bandpass profiles. Finally, the spectra were shifted to heliocentric velocity. In this paper we selected isolated spectral lines and measured their EWs to find spectral indicators of stellar physical parameters (\teff\ and \logg). However, our targets do not cover a sufficiently wide range of [Fe/H] to investigate the relation between EWs and metallicity. Therefore, the relations between EWs and stellar parameters (\teff\ and \logg) derived from their indicators are applicable exclusively for the stars with the near-solar abundances. We have found several spectral indicators of \teff\ such as Al I 1.672, Al I 2.117, Na I 2.209, Ti I 2.224, and CO 2.293~\um. Among the five spectral features, Al I 2.117 and Ti I 2.224~\um\ lines are newly found spectral indicators of \teff. Most of these spectral lines do not have significant correlations with \logg\ (Figure~\ref{ew_H_logg} and Figure~\ref{ew_K_logg}), but CO 2.293~\um\ EWs correlate with \teff\ and also \logg\ (Figure~\ref{co_ew}). Therefore, we found a relation among the EW of the CO feature, \teff, and \logg\ using a regression analysis (Equation~\ref{eq_regress}). Surface gravity can be calculated by measuring the EW of the CO overtone band 2.293~\um\ and the \teff\ derived from other lines (Table~\ref{tbl_diagnostics}) including the CO feature itself. The proposed method to estimate the stellar \teff\ and \logg\ relies on the features in the K-band that are less affected by reddening than shorter wavelength NIR and optical spectroscopy. In addition, the spectral indicators are close in wavelength, so they can be observed quickly and efficiently with spectrographs that have high-resolution, but shorter wavelength coverage than IGRINS, making this diagnostic method easily accessible. In conclusion, the derived relations (Table~\ref{tbl_diagnostics}) provide empirical ways to estimate the \teff\ and the \logg\ of late-type stars without direct comparison to stellar atmospheric models for all the sources, which is a valuable tool for population studies. \clearpage | 18 | 8 | 1808.06629 |
1808 | 1808.03710_arXiv.txt | We present a newly implemented single-pulse pipeline for the PALFA survey to efficiently identify single radio pulses from pulsars, Rotating Radio Transients (RRATs) and Fast Radio Bursts (FRBs). % We have conducted a sensitivity analysis of this new pipeline in which multiple single pulses with a wide range of parameters were injected into PALFA data sets and run through the pipeline. Based on the recovered pulses, we find that for pulse widths $\rm < 5\ ms$ the sensitivity of the PALFA pipeline is at most a factor of $\rm \sim 2$ less sensitive to single pulses than our theoretical predictions. For pulse widths $\rm > 10\ ms$, as the $\rm DM$ decreases, the degradation in sensitivity gets worse and can increase up to a factor of $\rm \sim 4.5$. Using this pipeline, we have thus far discovered 7 pulsars and 2 RRATs and identified 3 candidate RRATs and 1 candidate FRB. % The confirmed pulsars and RRATs have DMs ranging from 133 to 386 pc~cm$^{-3}$ and flux densities ranging from 20 to 160 mJy. The pulsar periods range from 0.4 to 2.1 s. We report on candidate FRB 141113, which we argue is likely astrophysical and extragalactic, having $\rm DM \simeq 400\ pc~cm^{-3}$, which represents an excess over the Galactic maximum along this line of sight of $\rm \sim$ 100 - 200~pc~cm$^{-3}$. We consider implications for the FRB population and show via simulations that if FRB 141113 is real and extragalactic, the slope $\alpha$ of the distribution of integral source counts as a function of flux density ($N (>S) \propto S^{-\alpha}$) is $1.4 \pm 0.5$ (95\% confidence range). However this conclusion is dependent on several assumptions that require verification. | \label{sec:intro} Pulsars are rapidly rotating, highly magnetized neutron stars (NSs). The majority of currently known pulsars are best detected through their time-averaged emission. Pulsar surveys like the PALFA survey \citep[Pulsar Arecibo L-band Feed Array;][]{cfl+06} generally use Fast Fourier Transform (FFT) searches in the frequency domain to search for pulsars. However, radio pulsar surveys often suffer from the presence of red noise generated by receiver gain instabilities and terrestrial interference. This can reduce sensitivity, particularly to long-period pulsars. For example, \citet{lbh+15} reported that due to the presence of red noise, % the sensitivity of the PALFA survey is significantly degraded for periods $\rm P > 0.5\ s$, with a greater degradation in sensitivity for longer spin periods. In order to mitigate such problems, more effective time domain searches like the fast-folding algorithm ~\citep[FFA, see][and references therein]{lk05,kml+09,pkr+18} and single-pulse search techniques ~\citep[as described by][]{cm03} can be used. Rotating Radio Transients (RRATs) are a relatively recently discovered class of NSs that were detected only through their individual pulses~\citep{mll+06}. % Due to the sporadic nature of their emission, surveys cannot rely on standard FFT searches to effectively look for RRAT signals. Instead, single-pulse search techniques % are required. Fast Radio Bursts (FRBs) are also a recently discovered phenomenon characterized by short (few $\rm ms$) radio bursts with high dispersion measures ($\rm DMs$)~\citep{lbm+07}. Unlike RRATs, which have observed $\rm DMs$ smaller than the maximum Galactic $\rm DM$ along the line of sight as predicted by Galactic free electron density models \citep{cl03,ymw16}, FRBs have $\rm DM$s that are much larger than this, implying extragalactic or even cosmological distances. To date, $\rm 34$ FRBs have been discovered\footnote{\url{www.frbcat.org}}, with only one FRB seen to repeat~\citep{ssh+16}. Like RRATs, FRBs can only be detected via single pulse-search techniques due to their transient nature. It is important to understand a survey's sensitivity to FRBs and RRATs as a function of various parameters (such as pulse width, $\rm DM$, scattering measure) if one is to accurately characterize the underlying sky event rates of these sources for population studies. The PALFA Survey is the most sensitive wide-area survey for radio pulsars and short radio transients ever conducted. Operating at a radio frequency band centered at $\rm 1.4\ GHz$, PALFA searches the Galactic plane ($\rm |b| < 5^{\circ}$), using the Arecibo Observatory, the 305-m single dish radio telescope located in Arecibo, Puerto Rico~\citep[see][for more details]{cfl+06,dcm+09,lbh+15}. Since the survey began in $\rm 2004$, it has discovered $\rm 178$ pulsars, including 15 RRATs and one FRB. ~\cite{lbh+15} comprehensively characterized the sensitivity of PALFA to radio pulsars, and showed that it is sensitive to millisecond pulsars as predicted by theoretical models based on the radiometer equation which assumes white noise. However, PALFA suffers significant degradation to long-period pulsars due to the presence of red noise in the data. In order to improve the search for long-period pulsars, the PALFA collaboration has introduced a fast-folding algorithm~\citep{pkr+18}. \cite{dcm+09} described an early single-pulse search algorithm for PALFA, reporting on the discovery of seven objects. % Here, we describe a new single-pulse search pipeline that we have also introduced to help identify long-period pulsars, RRATs and FRBs in our data. This new pipeline is described in \S\ref{sec:pipeline}. In \S\ref{sec:sensitivity}, we describe the survey's sensitivity to single pulses using an injection analysis. In \S\ref{sec:discoveries} we report new and candidate astrophysical sources discovered by this pipeline. We discuss a new candidate FRB, FRB 141113 in \S\ref{sec:frb}, and its implications for the FRB population in \S \ref{sec:implications}. We present our conclusions in \S\ref{sec:conclusion}. | \label{sec:conclusion} We have described a new, more systematic single-pulse pipeline to improve the search for pulsars, RRATs, and FRBs in the PALFA survey. The pipeline adds post-processing features to efficiently identify astrophysical single pulses. We also performed a robust sensitivity analysis of the PALFA survey to single pulses using injection of synthetic signals into survey data. % We find that for pulse widths $\rm < 5\ ms$ our survey is at most a factor of $\rm \sim 2$ less sensitive to single pulses than the theoretical predictions. For pulse widths $\rm > 10\ ms$, as the $\rm DM$ decreases, the degradation in sensitivity gets worse by up to a factor of $\rm \sim 4.5$. In order to better understand the actual sensitivities to single pulses in various radio transient surveys, we recommend similar characterization of their deployed detection pipelines. Using our pipeline, we have discovered one pulsar and two RRATs that were not detected using periodicity searching techniques, six pulsars that were detected by both single pulse and periodicity pipelines, three candidate RRATs, and one candidate FRB. This latter source, FRB 141113, has a DM more than twice the likely Galactic maximum along the line of sight, and multi-wavelength observations show it is very likely to be extragalactic. If so, it is consistent with being one of the lowest luminosity FRBs yet discovered. Simulations accounting for the sensitivity of PALFA and the discovery of FRB 121102 in addition to this new source indicate that the slope of the log $\rm N$--log $\rm S$ relation for the FRB population (i.e., $N(>S) \propto S^{-\alpha}$) is $\alpha = 1.4 \pm 0.5$ (95\% confidence). The steepness of that distribution is at odds with previous suggestions of a much flatter slope \citep{vedantham2016}. However, relaxing some reasonable assumptions in our calculation results in somewhat lower mean slopes, with uncertainty ranges that still bracket flatter population distributions. | 18 | 8 | 1808.03710 |
1808 | 1808.08258_arXiv.txt | We report the detection of methanol in the disk around the young outbursting star V883 Ori with the Atacama Large Millimeter/submillimeter Array (ALMA). Four transitions are observed with upper level energies ranging between 115 and 459 K. The emission is spatially resolved with the 0.14$\arcsec$ beam and follows the Keplerian rotation previously observed for C$^{18}$O. Using a rotational diagram analysis, we find a disk-averaged column density of $\sim10^{17}$ cm$^{-2}$ and a rotational temperature of \mbox{$\sim90-100$ K}, suggesting that the methanol has thermally desorbed from the dust grains. We derive outer radii between 120 and 140 AU for the different transitions, compared to the 360 AU outer radius for C$^{18}$O. Depending on the exact physical structure of the disk, the methanol emission could originate in the surface layers beyond the water snowline. Alternatively, the bulk of the methanol emission originates inside the water snowline, which can then be as far out as $\sim$100 AU, instead of 42 AU as was previously inferred from the continuum opacity. In addition, these results show that outbursting young stars like V883 Ori are good sources to study the ice composition of planet forming material through thermally desorbed complex molecules, which have proven to be hard to observe in more evolved protoplanetary disks. | \label{sec:introduction} Snowlines in disks around young stars mark the midplane locations where molecular species freeze out from the gas phase onto dust grains. The most important snowline for planet formation is the water snowline. Planetesimal formation is expected to be significantly enhanced at this snowline because the bulk of the ice mass is in water ice \citep[e.g.,][]{Stevenson1988,Ros2013,Schoonenberg2017a}. In addition, the elemental composition of both the gas and ice changes across this snowline since water is a major carrier of oxygen. The bulk composition of planets therefore depends on their formation location with respect to the water snowline \citep[e.g.,][]{Oberg2011,Madhusudhan2014,Eistrup2018}. \begin{table*}[ht!] \caption{Overview of the molecular line observations toward V883 Ori. \label{tab:Lineparameters}} \centering \begin{tabular}{l c c c c c c c c c} \hline\hline Species & Transition & Frequency & $A_{\rm{ul}}$ & $E_{\rm{up}}/k$ & $g_{\rm{up}}$ & $F_{\rm{peak}}$ & $F_{\rm{int}}$\tablenotemark{a} & $R_{\rm{out}}$ \\ & & (GHz) & (s$^{-1}$) & (K) & & (mJy beam$^{-1}$) & (Jy km s$^{-1}$) & (AU) \\ \hline C$^{18}$O & $2-1$ & 219.560354 & 6.03$\times$10$^{-7}$ & 16 & 3 & 117 $\pm$ 6 & 1.1 $\pm$ 0.08 & 361 $\pm$ 23 \\ CH$_3$OH (A$^-$) & $5_4-6_3$ & 346.202719 & 2.12$\times$10$^{-5}$ & 115 & 11 & \hspace{0.3cm}105 $\pm$ 21\tablenotemark{b} & \hspace{0.15cm}1.6 $\pm$ 0.18\tablenotemark{b} & \hspace{0.15cm}142 $\pm$ 27\tablenotemark{b} \\ CH$_3$OH (A$^+$) & $5_4-6_3$ & 346.204271 & 2.12$\times$10$^{-5}$ & 115 & 11 & \hspace{0.3cm}105 $\pm$ 21\tablenotemark{b} & \hspace{0.15cm}1.6 $\pm$ 0.18\tablenotemark{b} & \hspace{0.15cm}142 $\pm$ 27\tablenotemark{b} \\ CH$_3$OH (A$^-$) & $16_1-15_2$ & 345.903916 & 8.78$\times$10$^{-5}$ & 333 & 33 & \hspace{0.15cm}108 $\pm$ 21 & 1.1 $\pm$ 0.18 & 136 $\pm$ 18 \\ CH$_3$OH (E2) & $18_3-17_4$ & 345.919260 & 7.10$\times$10$^{-5}$ & 459 & 37 & \hspace{0.3cm}77 $\pm$ 21 & 0.5 $\pm$ 0.18 & 117 $\pm$ 26\\ \hline \end{tabular} \tablenotetext{a}{Within a 1.0$\arcsec$ aperture for C$^{18}$O and within a 0.6$\arcsec$ aperture for CH$_3$OH.} \tablenotetext{b}{Between 0.5 and 7.0 km s$^{-1}$, that is, for both lines combined.} \end{table*} For water, the transition from ice to gas occurs when the temperature exceeds roughly 100 K \citep{Fraser2001}. This places the snowline at a few AU from the star in protoplanetary disks, making it hard to observe. However, heating is temporarily enhanced during protostellar accretion bursts, which causes ices to sublimate out to larger radii. After such a burst, the circumstellar dust cools rapidly \citep{Johnstone2013}, while for the molecules it takes much longer to freeze back onto the grains \citep{Rodgers2003}. As a result snowlines are shifted away from the star \citep{Lee2007,Visser-Bergin2012,Vorobyov2013,Visser2015}. This has been observed for CO toward a sample of protostars \citep{Jorgensen2015,Frimann2017}. V883 Ori is a FU Orionis object in the Orion A L1641 molecular cloud ($d \sim 400$ pc; \citealt{Kounkel2017}) with a bolometric luminosity of $\sim218 L_{\sun}$ \citep{Strom1993,Furlan2016}. Although the onset of the V883 Ori outburst was not directly observed, evidence for an ongoing outburst that began before 1888 comes from its associated reflection nebula \citep{Pickering1890} and the similarity of its near-IR spectrum to that of FU Ori \citep{Connelley2018}. The 1.3 $M_{\sun}$ star is surrounded by a $\gtrsim$0.3 M$_{\sun}$ rotationally supported disk \citep{Cieza2016,Cieza2018} still embedded in its envelope. The location of the water snowline in V883 Ori was inferred from a change in the continuum opacity at 42 AU \citep{Cieza2016}, which may be due to a pileup of dust interior to the snowline \citep{Birnstiel2010,Banzatti2015,Pinilla2016}, or water evaporation and re-coagulation of silicates \citep{Schoonenberg2017b}. However, the origin of various structures seen in continuum emission of disks is still heavily debated and radial discontinuities in the spectral index are not necessarily related to snowlines \citep{vanTerwisga2018}. Molecular observations are thus needed to confirm or refute the snowline location. Unfortunately, water is hard to observe from the ground and warm water ($T \gtrsim 100$ K) has not yet been observed in young disks (Harsono et al., in prep.), so observing the snowline directly is difficult. A complementary approach is to observe other molecules whose distribution can be related to the snowline. Methanol (CH$_3$OH) is generally used to probe the $\gtrsim$ 100 K region in hot cores \citep{herbst2009}, because its volatility is similar to that of water \citep[e.g.,][]{Brown2007}. We serendipitously detected spatially resolved methanol emission in the V883 Ori disk with the VLA/ALMA Nascent Disk And Multiplicity (VANDAM) Orion survey that aims to characterize the embedded disks in Orion (PI: Tobin, Tobin et al., in prep.). Analysis of the methanol observations and comparison with earlier C$^{18}$O observations shows that the methanol is thermally desorbed and suggests that the water snowline can be as far out as $\sim100$ AU. \vspace{0.5cm} | \label{sec:discussion} \subsection{Location of the water snowline} The distribution of CH$_3$OH, and hence the relationship between its emission and the water snowline, depends on the physical structure of the disk, as illustrated in Figure~\ref{fig:Cartoon}. CH$_3$OH is present in the gas phase where the temperature exceeds the thermal desorption temperature of $\sim$100 K, and where there is a sufficiently large column of material to shield the UV radiation and prevent photodissociation ($A_V \geq 3$). In the surface layers, the photodissociation timescale is tens of years \citep{Heays2017}, comparable to the outburst duration (10-100 years). This means that the radial extent of the methanol layer higher up in the disk beyond the midplane water snowline is set by the intercept of the snow surface and the $A_V=3$ contour. In addition, the magnitude of the CH$_3$OH column density drop across the water snowline depends on the height of the snow surface; the higher up in the disk the larger the drop. Whether the emission then traces this column density profile depends on the optical depth of both the CH$_3$OH as well as the dust. \begin{figure*}[ht!] \centering \includegraphics[width=\textwidth,trim={0cm 18.6cm 0cm 1.5cm},clip]{Figure4.pdf} \caption{Illustration showing two possible scenarios for the distribution of CH$_3$OH with respect to the water snowline (\textit{top}), and the corresponding column density profiles (\textit{bottom}). Distinguishing between the scenarios \textbf{requires} a detailed physical model of the V883 Ori disk. The region where CH$_3$OH is present in the gas phase is highlighted in red. If the CH$_3$OH emission is optically thick, only molecules above the $\tau=1$ surface (red line) are observed. The location of the $\tau=1$ surface depends on both the disk structure and the CH$_3$OH transition. The corresponding column density traced in the optically thick case is shown with a dashed line, while the solid line represents the total CH$_3$OH column. The blue curve indicates the water snow surface at $\sim$100 K, and the vertical blue line the corresponding midplane water snowline. CH$_3$OH desorbs at a similar temperature as water. The dashed black line marks the $A_V=3$ contour above which CH$_3$OH is photodissociated. The C$^{18}$O column density profile is shown in orange.} \label{fig:Cartoon} \end{figure*} Due to this interplay of several parameters, a detailed physical model of the disk is required to derive the water snowline location from methanol emission. It may thus be possible to see CH$_3$OH emission out to $\sim120-140$AU while the snowline is around 40 AU. This would require for example a water snow surface close to the midplane, and/or optically thick methanol emission (Figure~\ref{fig:Cartoon}, left panel). The vertical temperature structure is strongly dependent on the dust distribution and settling of the large grains can result in a steep vertical temperature profile with the snow surface closer to the midplane than in the case of less grain settling \citep[e.g.,][]{Facchini2017}. However, especially if the emission is optically thin, the bulk of the methanol emission is more likely to originate inside the water snowline (Figure~\ref{fig:Cartoon}, right panel), since the CH$_3$OH column density can drop $\sim$3 orders of magnitude crossing the snowline assuming a constant abundance for the gas-phase CH$_3$OH \citep[see e.g., the simple model for the CO snowline in][]{Qi2013}. Assuming a step function for the column density, this would mean that the snowline in V883 Ori can be as far out as $\sim100-125$ AU, taking into account the 40 AU beam by deconvolving the radial profiles. Non-thermal desorption processes are not expected to influence the relationship between CH$_3$OH and the water snowline. Such processes have been invoked to explain the CH$_3$OH emission in TW Hya \citep{Walsh2016}, but this required gas-phase CH$_3$OH outside the CO snowline ($T\lesssim$ 20~K) at an abundance of $\sim10^{-12}-10^{-11}$, several orders of magnitude lower than observed here and expected from ice abundances (Sect.~\ref{subsec:rd}). Besides detailed modeling, observations of other molecular tracers could put better constraints on the water snowline location. H$^{13}$CO$^+$ has shown to be a promising tracer in the envelope around NGC1333 IRAS2A \citep{vantHoff2018}, because the main destroyer of HCO$^+$ is gas-phase water. \vspace{1.0cm} \subsection{Ice composition of planet forming material} One of the key questions in planet formation is whether planetary systems inherit their chemical composition from the natal cloud or whether the material is significantly processed en route to the disk. Observations of many complex molecules, including methanol, around young protostars at solar system scales \citep[e.g.,][]{Jorgensen2016} and in comets \citep[e.g.,][]{Mumma2011,LeRoy2015} show that a large complexity is present during both the early as well as the final stages of planet formation. The chemical complexity in protoplanetary disks, however, is hard to probe. Due to the low temperatures ($<$100 K), complex molecules are frozen out onto dust grains at radii larger than a few AU, and ices can only be observed through infrared absorption in edge-on systems. Although alternative desorption processes may get these molecules into the gas phase, as has been shown for water \citep{Hogerheijde2011}, so far only CH$_3$OH and CH$_3$CN have been observed in disks \citep{Oberg2015b,Walsh2016,Bergner2018,Loomis2018}. Moreover, as it is unclear which processes operate for which species and what the efficiencies are, the observed gas composition cannot directly be linked to the ice composition of planet forming bodies. The results presented here show that complex molecules can thermally desorb in disks around young stars that have recently undergone an accretion burst. Moreover, their emission extends out to more than 100 AU around V883 Ori and is readily detected and spatially resolved with only one minute of integration with ALMA. V883 Ori is the longest-lasting known outburst and one of, if not the, most luminous. This makes it an archetype for understanding disk chemistry in fainter outbursts. It also provides a look at how younger outbursts may evolve over the century following their outbursts. Young disks like V883 Ori thus provide the unique opportunity to study the chemical complexity at the onset of planet formation. | 18 | 8 | 1808.08258 |
1808 | 1808.10790_arXiv.txt | {In 2015, Barman et al. presented detections of absorption from water, carbon monoxide, and methane in the atmosphere of the directly imaged exoplanet HR8799b using integral field spectroscopy (IFS) with OSIRIS on the Keck II telescope. We recently devised a new method to analyse IFU data, called {\sl molecule mapping}, searching for high-frequency signatures of particular molecules in an IFU data cube.} {The aim of this paper is to use the molecule mapping technique to search for the previously detected spectral signatures in HR8799b using the same data, allowing a comparison of molecule mapping with previous methods.} {The medium-resolution H and K-band pipeline-reduced archival data were retrieved from the Keck archive facility. Telluric and stellar lines were removed from each spectrum in the data cube, after which the residuals were cross-correlated with model spectra of carbon monoxide, water and methane.} {Both carbon monoxide and water are clearly detected at high signal-to-noise, however, methane is not retrieved.}{Molecule mapping works very well on the OSIRIS data of exoplanet HR8799b. However, it is not evident why methane is detected in the original analysis, but not with the molecule mapping technique. Possible causes could be the presence of telluric residuals, different spectral filtering techniques, or the use of different methane models. We do note that in the original analysis methane was only detected in the K-band, while the H-band methane signal could be expected to be comparably strong. More sensitive observations with the JWST will be capable of confirming or disproving the presence of methane in this planet at high confidence.} | The planetary system HR8799 is unique in that it hosts four directly-imaged planets \citep{Marois2008,Marois2010}. The A-star is estimated to have an age of 30 Myr \citep{Marois2010} and is located at a distance of 39 pc \citep{Marois2008}. The planets orbit at 15 to 70 au between a warm dust belt and a cold Kuiper belt-like structure \citep{Su2009, Booth2016}, implying orbital periods between 50 and 500 years. Indeed, recent time series observations clearly show the orbital motion of the planets \footnote{http://jasonwang.space/orbits.html}. HR8799b is the outermost planet, 1.7 arcseconds away from the host star, with a K-band magnitude of Ks=14.05, corresponding to a contrast of $\sim3300$ at these wavelengths \citep{Marois2008}. Comparisons of the planet's spectral energy distribution with planet evolutionary models provide a mass estimate of $M=5^{+2}_{-1} M_{\rm{J}}$ with a planet radius of $R=1.2\pm0.1 R_{\rm{J}}$ \citep{Currie2011,Soummer2011,Wright2011}. Early observations with the integral field spectrograph OSIRIS \citep{OSIRIS2003,OSIRIS2006} on the Keck II telescope at $2.1 - 2.2\ \mu $m combined with flux measurements at other wavelengths show that the near-infrared spectrum of HR8799b is consistent with that of an early T-dwarf with an intermediate cloud deck and possibly a weak methane feature \citep{Bowler2010}. Atmospheric modeling by \citet{Madhusudhan2011} infers a thick cloud cover, and shows that the 3.3 $\mu $m flux of the planet is inconsistent with the expected fiducial methane abundance - requiring significantly less methane (see also \cite{Marley2012,Skemer2012}). Further observations by \citet{Oppenheimer2013} and \citet{Ingraham2014} also show no significant methane absorption. However, later observations with OSIRIS at H and K-band do provide evidence for methane \citep{Barman2011,Barman2015}. A binned-down, low resolution spectral analysis of the IFU data \citep{Barman2011} shows strong water absorption in H and K, and weak methane and CO in K-band. Subsequent analysis at the full resolving power of R=4000 \citep{Barman2015} using cross-correlation techniques, seem to confirm the methane and carbon monoxide detections in K-band, but show no evidence for methane absorption in the H-band data, while this is expected from the absorption feature at 1.65 $\mu$m. In this paper, we apply a different analysis on the existing OSIRIS H and K-band data, also making use of cross-correlation techniques. The technique, dubbed {\it molecule mapping}, targets the high-frequency signatures of particular molecules in an IFU data cube, ultimately providing a measure of the amount of evidence for molecular absorption at all spatial locations in the IFU data cube \citep{Hoeijmakers2018}. The technique is described in detail in \cite{Hoeijmakers2018}, where it is applied to SINFONI/VLT data of the beta Pictoris system, showing strong absorption from both CO and H$_2$O at SNRs$\sim$15 at the location of exoplanet beta Pictoris b. First, a stellar reference spectrum is constructed using the spaxels near the star position. Subsequently, a scaled version of this stellar reference is subtracted from the spectra at each position in the IFU field of view. The redisual spectra are then cross-correlated with molecular templates. The cross-correlation co-adds the individual absorption lines of the planet emission spectrum constructively, while this is not the case for (residual) telluric and stellar features \citep{Hoeijmakers2018}. The method depends neither on field rotation, as is the case for Angular Differential Imaging (ADI; \cite{Marois2006}) methods, nor on wavelength scaling of the PSF, as for Spectral Differential Imaging (SDI; \cite{Mawet2012}) techniques, making it particularly powerful at small angular distances from the host star where residual speckles are strongest. Molecule mapping is different from the technique utilized by \cite{Barman2015}. In particular, they perform telluric calibrations using an early-type standard star, which implicitly also removes possible stellar absorption lines. The observations are described in Section \ref{sec:observations}, and the data analysis is presented in Section \ref{sec:analysis}. The resulting molecule maps are shown and discussed in Section \ref{sec:results}. | \label{sec:results} \subsection{Molecule maps} The resulting molecule maps consist of the cross-correlation signal from each spaxel at the planet radial velocity, here assumed to be 35 km sec$^{-1}$ (blueshifted) for all nights, originating from the combination of the barycentric velocity of the observatory (similar for all observations) and the system velocity of the star. The resolution of the cross-correlation function in both bands is also about 35 km sec$^{-1}$. Figure \ref{fig:mmH} presents the molecule maps from the H-band data, showing water and methane respectively. While water is marginally detected at SNR$\sim$2, there is no sign of methane. The SNRs are determined as the peak signal at the planet velocity divided by the standard devation of the cross-correlation function over a wide range of velocity ($\pm15000$ km sec$^{-1}$). Figure \ref{fig:mmK} presents the molecule maps from the K-band data, showing from top-left to bottom-right water, carbon monoxide, methane, and ammonia. Both carbon monoxide and water are clearly detected at SNR$\sim$6, while again there is no sign of methane, and neither of ammonia. \begin{figure} \centering \includegraphics[scale=.5]{Figure2.pdf} \caption{\label{fig:mmH}Molecule maps from the H-band data showing water on the left and methane on the right. The maps are centered on the location of the planet. The planet is detected in the map of water, but not in that of methane.} \end{figure} \begin{figure} \centering \includegraphics[scale=.5]{Figure3.pdf} \caption{\label{fig:mmK} Molecule maps from the K-band data showing in the top row left to right carbon monoxide and water and in the bottom row methane and ammonia. The maps are centered on the location of the planet. The planet is detected in the maps of water and carbon monoxide, but not in the maps of methane and ammonia.} \end{figure} The top panel of Figure \ref{fig:cc-H-results} shows the normalized and filtered water and methane model spectra used for the cross-correlation in H-band, and the processed spectrum at the planet position. The lower panel shows the cross-correlation functions for these molecules at this planet position, showing the low-confidence signal for water, and the absence of a methane signal. Figure \ref{fig:cc-K-results} shows the same for the K-band data, including carbon monoxide, water, methane and ammonia. The cross correlation functions show the clear detection for water and carbon monoxide at the planet velocity, but no sign of methane and ammonia. The resulting signal-to-noise ratios are summarized in Table \ref{tab:SNR}. \begin{figure} \centering \resizebox{\hsize}{!}{\includegraphics{cc_H.jpeg}} \caption{\label{fig:cc-H-results} \emph{Top:} Filtered spectrum of HR8799b (black) and template spectra of water (blue) and methane (red) in the H band, offset by an arbitrary amount. \emph{Bottom:} Cross-correlation functions of the planet spectrum with the model spectra plotted above.} \end{figure} \begin{figure} \centering \resizebox{\hsize}{!}{\includegraphics{cc_K.jpeg}} \caption{\label{fig:cc-K-results} As Figure \ref{fig:cc-H-results}, but in the K band and with the addition of methane (green) and ammonia (cyan) model spectra and cross-correlation functions. The difference in the resolution of the CCF is due to the difference in resolution of the H and K band data.} \end{figure} \begin{table}[h!] \centering \begin{tabular}{|c|c|c|} \hline Molecule & H band & K band\\ \hline $H_2O$ & 2.2 & 6.2 \\ $CH_4$ & 0.4 & 0.8 \\ $CO$ & - & 6.4 \\ $NH_3$ & - & 0.6 \\ \hline \end{tabular} \caption{\label{tab:SNR}Signal to noise ratios of the cross-correlation functions in the H and K bands at a the planet radial velocity of -35 km sec$^{-1}$. } \end{table} No significant molecule signals are found away from the planet position. This means that the telluric line removal was successful, and that possible telluric residuals have not resulted in spurious signals. \subsection{Comparison to earlier work} We directly compare our results with those obtained by \cite{Barman2011} and \cite{Barman2015}, who used the same OSIRIS data - albeit we use a total of 50\% and 40\% fewer on-target exposures than the previous analysis due to the availability of reduced data cubes in the Keck archive. An equally powerful analysis would result in $\sim$30\% lower SNRs. Although \cite{Barman2011} and \cite{Barman2015} do not explicitly quote a statistical significance for their detections, visual inspection of their cross-correlation functions indicate that our detections of H$_2$O and CO, in particular in the K-band, are significantly stronger. We also note that the velocity resolution of our cross-correlation functions are higher. This may in part explain why our analysis results in a higher confidence for these molecules.s In contrast, while \cite{Barman2015} present a significant detection of methane in the K-band OSIRIS data, we see no sign of this molecule in our analysis. Possible causes could be the presence of telluric residuals in either of the analyses, the differences in spectral filtering techniques, or the use of different methane models. In particular, they perform telluric calibrations using an early-type standard star, while our analysis performs effectively a 'self-calibration' for telluric line removal using the data themselves. We also note that in the original analysis methane was only detected in the K-band, while the H-band methane signal could be expected to be comparably strong. Future more sensitive observations, e.g. with the JWST, should reveal whether methane is really present in the atmosphere of HR8799b or not. | 18 | 8 | 1808.10790 |
1808 | 1808.07596_arXiv.txt | Identification of the cosmic-ray (CR) ``PeVatrons'', which are sources capable of accelerating particles to $\sim10^{15}$~eV energies and higher, may lead to resolving the long-standing question of the origin of the spectral feature in the all-particle CR spectrum known as the ``knee''. Because CRs with these energies are deflected by interstellar magnetic fields identification of individual sources and determination of their spectral characteristics is more likely via very high energy \gray{} emissions, which provide the necessary directional information. However, pair production on the interstellar radiation field (ISRF) and cosmic microwave background (CMB) leads to steepening of the high-energy tails of \gray{} spectra, and should be corrected for to enable true properties of the spectrum at source to be recovered. Employing recently developed three-dimensional ISRF models this paper quantifies the pair-absorption effect on spectra for sources in the Galactic centre (GC) direction at 8.5~kpc and 23.5~kpc distance, with the latter corresponding to the far side of the Galactic stellar disc where it is expected that discrimination of spectral features $>10$~TeV will be possible by the forthcoming Cherenkov Telescope Array (CTA). The estimates made suggest spectral cutoffs could be underestimated by factors of a few in the energy range so far sampled by TeV \gray{} telescopes. As an example to illustrate this, the recent HESS measurements of diffuse \gray{} emissions possibly associated with injection of CRs nearby Sgr~A$^*$ are ISRF-corrected, and estimates of the spectral cutoff are re-evaluated. It is found that it could be higher by up to a factor $\sim 2$, indicating that these emissions may be consistent with a CR accelerator with a spectral cutoff of at least 1\,PeV at the 95\% confidence level. | 18 | 8 | 1808.07596 |
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1808 | 1808.09632_arXiv.txt | \label{sec:intro} Exoplanet direct imaging is necessarily the future of exoplanet science. RV surveys have a very important role to play in discovering exoplanets and measuring their mass, but as an indirect detection method its science return has limited scope. Transit surveys help to fill in this information, but are hampered by the small fraction of systems that transit and practical limits on atmospheric spectroscopy for atmospheres with small scale heights, especially for planets with long periods. Direct sensing at high contrast allows us to directly image and obtain spectra of planets. Despite the technique's obvious attractions, it is technically demanding. The first generation of dedicated facility planet imagers, GPI and SPHERE, were optimistic in their assumptions about the achieved contrast and exoplanet populations at wide separations, and as with many first attempts encountered unforeseen issues. Despite this, SPHERE has met its design contrast goals, and GPI should meet or exceed them with investment. New instrumentation for 30-m telescopes (GSMTs, Giant Segmented Mirror Telescopes) will benefit from 10--20 years of progress in our understanding of direct imaging, significantly improved technology, a much clearer picture of the underlying demographics of exoplanets, and a decade of experimentation on testbeds like SCExAO and MagAO-X to improve contrast. Direct imaging is beginning to climb the power law of the exoplanet distribution. Given funding and the space to innovate, scientific progress will be rapid. | 18 | 8 | 1808.09632 |
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1808 | 1808.10103_arXiv.txt | The SOL2014-09-01 far-side solar eruptive event produced hard electromagnetic and radio emissions observed with detectors at near-Earth vantage points. Especially challenging was a long-duration $>100$\,MeV $\gamma$-ray burst probably produced by accelerated protons exceeding 300\,MeV. This observation raised a question of how high-energy protons could reach the Earth-facing solar surface. Some preceding studies discussed a scenario in which protons accelerated by a CME-driven shock high in the corona return to the solar surface. We continue with the analysis of this challenging event, involving radio images from the \textit{Nan{\c c}ay Radioheliograph} and hard X-ray data from the \textit{High Energy Neutron Detector} (HEND) of the \textit{Gamma-Ray Spectrometer} onboard the \textit{Mars Odyssey} space observatory located near Mars. HEND recorded unocculted flare emission. The results indicate that the emissions observed from the Earth's direction were generated by flare-accelerated electrons and protons trapped in static long coronal loops. Their reacceleration is possible in these loops by a shock wave, which was excited by the eruption, being initially not CME-driven. The results highlight the ways to address remaining questions. | \label{S-introduction} The source of solar energetic particles (SEPs) produced in solar eruptive-flare events is a subject of long-standing debate. SEPs consist of different species dominated by protons. Two sources of accelerated protons have been considered (see, \textit{e.g.}, \citealp{Kahler2001, Kallenrode2003, Aschwanden2012, Reames2013}). One presumable origin of accelerated protons is associated with flare processes in solar active regions manifested in X-rays and microwaves. Another source is related to a bow-shock driven by the outer surface of a super-Alfv{\'e}nic coronal mass ejection (CME). Many indications have been considered to identify the elusive source of accelerated protons. One of them is $\gamma$-ray emission, which was mostly observed concurrently with other flare emissions and seemingly favored the acceleration of protons in flares along with electrons (\textit{e.g.} \citealp{RamatyMandzhavidze2000, LivshitsBelov2004, ChuppRyan2009, Kurt2010, Vilmer2011}). Flare emissions are observed in a wide electromagnetic range, from radio waves up to high-energy $\gamma$-rays. Gyrosynchrotron emission observed in the radio range and a broadband hard X-ray (HXR) and $\gamma$-ray bremsstrahlung continuum are produced by accelerated electrons. Accelerated protons and heavier ions can be recognized from discrete $\gamma$-ray lines. Of special interest is the $\pi^0$-decay emission. Neutral pions appear in proton--proton collisions, when the proton energy exceeds 300\,MeV, and they rapidly decay into two photons, producing a Doppler-broadened wide enhancement in the $\gamma$-ray spectrum around 100\,MeV on top of the bremsstrahlung continuum (\textit{e.g.} \citealp{RamatyKozlovskyLingenfelter1975, HudsonRyan1995, Vilmer2011}). Thus, the $\pi^0$-decay emission is a direct indication of protons accelerated to high energies. Observations and identification of this $\gamma$-ray emission are only possible with a high sensitivity and sufficient spectral measurements at high energies. For this reason, the total number of events with a confident identification of the $\pi^0$-decay emission was fewer than 20 in the past (\textit{e.g.} \citealp{Ryan2000, Grechnev2008, ChuppRyan2009, Kurt2010, Kuznetsov2011}). Being temporally close to flare emissions produced by accelerated electrons, discrete nuclear $\gamma$-ray lines, and especially the $\pi^0$-decay emission, have been considered as evidence of proton acceleration in flares. On the other hand, $\gamma$-ray emission much longer than the HXR burst was observed in a few events (\textit{e.g.} \citealp{Forrest1985, Akimov1996, Ryan2000}). A challenge to the flare-related origin of $\gamma$-ray emission was provided by the observation of $\gamma$-ray emission from an event behind the solar limb. To explain this phenomenon, \cite{Cliver1993} proposed that protons accelerated by a CME-driven shock wave on an open magnetic field partly escaped into interplanetary space and partly returned to the solar surface, precipitating far from the flare region. With the advent of the \textit{Fermi Gamma-Ray Space Telescope} mission in 2008, high-sensitivity $\gamma$-ray observations become available with a comprehensive spectral information and coordinate measurements of $\gamma$-ray photons at $> 100$\,MeV by the \textit{Large Area Telescope} (LAT: \citealp{Atwood2009}). Although it is a non-solar mission, \textit{Fermi} also provides rich information for solar studies. \textit{Fermi} has shown that $\gamma$-ray emissions are quite common in solar flares. Thirty long-duration $\gamma$-ray events have been observed \citep{Share2017}. \cite{Pesce-Rollins2015} reported on the detection by \textit{Fermi}/LAT of high-energy $\gamma$-ray emissions from three behind-the-limb solar flares on 11 October 2013, 6 January 2014, and 1 September 2014. These events were addressed by \cite{Ackermann2017}. The $\pi^0$-decay emission was identified with confidence in two of them, SOL2013-10-11 and SOL2014-09-01. The authors invoked the idea of \cite{Cliver1993} to interpret these emissions. \cite{Plotnikov2017} elaborated on this idea in their analysis of the three events. Among the issues analyzed, by means of three-dimensional reconstructions of coronal shock fronts, the authors showed the events' magnetic connectivity to the Earth-facing solar surface. \cite{Ackermann2017} and \cite{Plotnikov2017} considered coronal shocks to be driven by fast CMEs and emphasized that the CME and associated shock wave were fastest in the 1 September 2014 event, where the high-energy $\gamma$-ray emission was strongest. On the other hand, \cite{Hudson2017} pointed out basic problems of the scenario proposed by \cite{Cliver1993}. First, a large mirror ratio at the base of an open coronal structure prevents the back-precipitation of particles from large coronal heights, so that only a small part of the protons is able to return to the Sun in this scheme. Second, the total number of high-energy protons estimated for a set of SEP events appears to be insufficient to sustain the high-energy $\gamma$-rays in the events addressed by \cite{Ackermann2017}. To explain the long-duration $\gamma$-rays from occulted events, \cite{Hudson2017} considered two options. In the ``Lasso'' scenario, some SEPs are captured in a noose, which extends to several solar radii and then retracts. In this scenario, trapped particles acquire energy due to the betatron acceleration and first-order Fermi process. The second option that he proposed is a ``coronal thick target'' scenario, in which protons trapped in a static volume generate pions and $\gamma$-ray continuum. As \cite{Hudson2017} estimated, this can proceed for a few hours. Analyzing the dynamic evolution of the global magnetic field and the shock wave considered to be CME-driven, \cite{Jin2018} simulated the CME in the 1 September 2014 event by using a global MHD model. The authors concluded that particles responsible for the high-energy $\gamma$-ray emission were accelerated in the CME environment and escaped the shock downstream region along magnetic fields connected to the solar surface far away from the flaring region. So one has to conclude that, in spite of the rich observational data available at present and a lot of efforts applied, the source of accelerated protons escapes identification. Furthermore, examining the ``flare \textit{vs.} CME-driven shock'' alternative, the researchers base their considerations on a simplified traditional scheme of the bow-shock excitation by the outer surface of a fast CME. However, recent studies show that coronal shock waves are initially excited by sharply erupting flux-ropes inside the developing CMEs, while reconnection processes underneath produce a flare (see, \textit{e.g.}, \citealp{Grechnev2016, Grechnev2018} for details and review). The flare, CME, and shock-formation processes turn out to be tightly associated, which determines a close relation between the characteristics of flares, CMEs, and shock waves. The situation gets still more complicated. If flare processes are actually responsible for acceleration of protons, then parameters of CMEs are misleading. If shock-acceleration is at work, then the acceleration starts earlier and at lower altitudes than assumed previously. If both sources are implicated, then untangling their contributions is still more difficult. Keeping in mind these circumstances, we continue with the analysis of the 1 September 2014 event. It was observed from different vantage points. From the Earth's direction, the HXR burst was observed by the \textit{Fermi Gamma-ray Burst Monitor} (GBM: \citealp{Meegan2009}) and the \textit{Wind}/Konus \textit{Gamma-Ray Burst Experiment} \citep{Aptekar1995}. A radio burst dominated by the gyrosynchrotron (GS) emission at frequencies $> 300$\,MHz was recorded by the \textit{Radio Solar Telescope Network} (RSTN: \citealp{Guidice1979, Guidice1981}), while its source was observed at the \textit{Nan{\c c}ay Radioheliograph} (NRH: \citealp{Kerdraon1997}). The GS burst was considered by \cite{Ackermann2017} and \cite{Carley2017}. The unocculted flare emission was recorded from the Martian direction by the \textit{High Energy Neutron Detector} (HEND) of the \textit{Gamma-Ray Spectrometer} onboard the \textit{Mars Odyssey} space observatory \citep{Boynton2004}. The SOL2014-09-01 event was not listed in the HEND catalog \citep{Livshits2017}; nevertheless, HEND actually observed it. Based on these data, we analyze the electromagnetic emissions observed, endeavor to figure out their possible sources, try to understand the causes of the long-lasting emissions, and reveal the history and possible role of the shock wave. In this way, we pursue understanding which scenarios of those proposed match the observations, specifying and refining some of the results and conclusions drawn previously. Section~\ref{S-emissions_sources} addresses electromagnetic emissions observed in the event and their probable sources. Section~\ref{S-shock_wave} analyzes shock waves and their kinematics. Section~\ref{S-discussion} discusses the results and their interpretation. Section~\ref{S-conclusion} summarizes the findings and presents the conclusion. | \label{S-conclusion} A combined analysis of observations of the far-side SOL2014-09-01 event from different vantage points has revealed the following circumstances. \begin{enumerate}[i] \item The lift-off of a hot (about 10\,MK) blob has been detected, which probably was an erupting flux rope. The blob rose radially and became the CME core. \item Unocculted flare emission consisted of two HXR peaks with similar spectra separated by 2.5 minutes. \item Each of the two flare peaks was preceded by the appearance of a shock wave by two to three minutes. \item The first HXR peak was manifested in a static off-limb gyrosynchrotron radio source of a corresponding duration and spectrum. \item The second HXR peak gave rise to a different static off-limb gyrosynchrotron radio source of a considerably longer duration and harder spectrum. This radio source was located in a system of long loops. \item The long-duration gyrosynchrotron burst from the second source was almost identical in shape with the HXR burst observed from the Earth's direction and rather similar to the $> 100$\,MeV $\gamma$-ray burst. All of these emissions could be produced by populations of electrons and protons injected into the long loops during the second flare burst and trapped there for a long time. \item The harder spectrum of the long-duration burst relative to the injection could be due to reacceleration of the particles trapped in closed loops by the second shock wave. \item The observations indicate that the sources of the gyrosynchrotron, HXR, and $\gamma$-ray emissions had a common location. It was considerably displaced with respect to the $> 100$\,MeV $\gamma$-ray emission centroid position. A probable key to the discrepancy is a contribution of $\gamma$-rays coming from high coronal structures and possibly the CME. The role of non-solar high-energy cosmic rays is not excluded. \end{enumerate} These findings can be reconciled within the following scenario. Two sharp eruptions probably occurred in Active Region 12158 with an interval of about 2.5 minutes. Each eruption impulsively excited a blast-wave-like shock, on the one hand, and resulted in strong particle acceleration in the flare site, on the other hand. Manifestations of the first flare peak were observed from the Earth's direction as an impulsive brightening of the arcade top. During the second peak, accelerated electrons and protons were injected into long loops, where they become trapped for a long time. The second shock wave possibly hit these loops obliquely, which resulted in reacceleration of trapped flare-accelerated electrons and protons. This presumable episode was not crucial; the long-duration gyrosynchrotron, hard X-ray, and $\gamma$-ray emissions were radiated from trapped particles, while reacceleration hardened their spectra. A presumable scenario with a shock-acceleration of particles high in the corona and their return to the solar surface along open magnetic structures meets basic difficulties and is not confirmed by observations. The region of trapped electrons and protons was located above the limb. Its connection to the Earth-facing solar surface near the Equator is not excluded, but does not seem to be necessary. While our analysis sheds additional light on this event, a number of issues remain to be addressed. We hope that our results will highlight possible ways for future studies. \begin{acks} This work is dedicated to the memory of M.A.~Livshits, who initiated this study. We appreciate discussions with E.~Carley, N.~Vilmer, and H.~Hudson, and useful remarks of the anonymous reviewer. We thank the NASA/SDO and the AIA and HMI science teams; the NASA's STEREO/SECCHI science and instrument teams; the teams of the SWAP telescope on the ESA's PROBA~2 spacecraft, the NASA's \textit{Fermi Gamma-Ray Space Telescope}, the \textit{Wind}/Konus team at the Ioffe Institute, the USAF RSTN network, and LASCO on SOHO. SOHO is a project of international cooperation between ESA and NASA. We thank the team maintaining the CME Catalog at the CDAW Data Center by NASA and the Catholic University of America in cooperation with the Naval Research Laboratory. The studies presented in Sections \ref{S-introduction}, \ref{S-overview}, and \ref{S-time_profiles} were carried out by V.~Kiselev and I.~Grigorieva and funded by the Russian Foundation of Basic Research under grant 17-32-50040\_mol\_nr. V.~Grechnev, A.~Kochanov, and A.~Uralov (Sections \ref{S-shock_wave}, \ref{S-discussion}, and \ref{S-conclusion}) were funded by the Russian Science Foundation under grant 18-12-00172. \end{acks} | 18 | 8 | 1808.10103 |
1808 | 1808.02534_arXiv.txt | Empirical models of supernova (SN) spectral energy distributions (SEDs) are widely used for SN survey simulations and photometric classifications. The existing library of SED models has excellent optical templates but limited, poorly constrained coverage of ultraviolet (UV) and infrared (IR) wavelengths. However, both regimes are critical for the design and operation of future SN surveys, particularly at IR wavelengths that will be accessible with the James Webb Space Telescope (JWST) and the Wide-Field Infrared Survey Telescope (WFIRST). We create a public repository of improved empirical SED templates using a sampling of Type Ia and core-collapse (CC) photometric light curves to extend the Type Ia parameterized SALT2 model and a set of SN Ib, SN Ic, and SN II SED templates into the UV and near-IR. We apply this new repository of extrapolated SN SED models to examine how future surveys can discriminate between CC and Type Ia SNe at UV and IR wavelengths, and present an open-source software package written in Python, \textit{SNSEDextend}, that enables a user to generate their own extrapolated SEDs. | } \newcommand{\data}{ \label{sub:intro} Photometric templates and models of supernova (SN) spectral energy distributions (SEDs) are critical tools for gleaning physical properties of supernovae (SNe) from observations, determining how those properties evolve over time, and performing SN classifications. Many SN analysis tools, such as the widely-used SuperNova ANAlysis \citep[SNANA][]{Kessler:2009a} and SNCosmo \citep{Barbary:2014} software packages, utilize a common set of empirically-derived SEDs that represent a variety of core-collapse (CC) and Type Ia SNe. Most existing template SEDs, however, are only constrained by data at optical wavelengths \citep{Blondin:2007}. Many of the software packages for photometric SN classification rely on these template SEDs \citep{Kessler:2010,Sako:2011}. Extending the SED templates into near-infrared (NIR) wavelengths is necessary for those classification tools to be applicable for the next generation of telescopes---such as the James Web Space Telescope (JWST), the Wide-Field Infrared Survey Telescope (WFIRST), and the Large Synoptic Survey Telescope (LSST)---which will provide a plethora of new SN observations that span wavelengths from the optical to the far-IR \citep{Dahlen:1999,Mesinger:2006,Ivezic:2008,Spergel:2015}. A preliminary extension of the SN SED library into NIR bands \citep{Pierel:2018a}\footnote{\href{dx.doi.org/10.5281/zenodo.1250492}{DOI:10.5281/zenodo.1250492}} has already been used for the analysis of SN discoveries in the Cosmic Assembly Near-infrared Deep Extragalactic Legacy Survey \citep[CANDELS,][]{Rodney:2014} and the Cluster Lensing And Supernova survey with Hubble \citep[CLASH,][]{Graur:2014}. The simplistic modified SN SED templates employed for that work have also been used to explore various survey strategies for the WFIRST SN program \citep{Hounsell:2017}. In this work we provide a more rigorous extension of ultraviolet (UV) and NIR coverage for current SEDs. First, we describe a new open-source software tool, \textit{SNSEDextend}, that is capable of extrapolating SN SEDs to match photometric observations. The data and methodology for extending CC SN SEDs are presented in Section~\ref{sec:ccsn}. Extrapolation of the Type Ia model SALT2 \citep{Guy:2010} is described in Section~\ref{sec:typeIa}. We then provide a new repository of SEDs extrapolated to cover the wavelength range $\sim 1700$--25,000~\mbox{\normalfont\AA}, and in Section~\ref{sec:results} we apply these SEDs to explore photometric SN classifications in IR bands. As the number of UV and IR observations of SNe increases, the accuracy of the extrapolations will continue to improve and the \textit{SNSEDextend} package will be available to supplement the repository with updated and new SED templates. Meanwhile, the intention is that these extrapolated SEDs will be used by the wider SN research community for simulations and photometric classifications. We note, however, that our extrapolations of the SALT2 Type Ia SN model to UV and NIR wavelengths are not intended to make SALT2 capable of light-curve fitting in those wavelength regimes for cosmological distance measurements. That would require retraining of the model, which is beyond the scope of this work. \ | } \newcolumntype{L}{>{$}l<{$}} % \newcolumntype{C}{>{$}c<{$}} % \newcolumntype{R}{>{$}r<{$}} % \begin{document} \title{Extending Supernova Spectral Templates for Next-Generation Space Telescope Observations} \newcommand{\USC}{Department of Physics and Astronomy, University of South Carolina, 712 Main St., Columbia, SC 29208, USA} \newcommand{\CfA}{Harvard-Smithsonian Center for Astrophysics, Cambridge, MA 02138, USA} \newcommand{\Chicago}{Department of Physics, The University of Chicago, Chicago, IL 60637, USA} \newcommand{\NYU}{Center for Cosmology and Particle Physics, New York University, New York, NY 10003, USA} \newcommand{\UCSB}{Department of Physics, University of California, Santa Barbara, CA 93106-9530, USA} \newcommand{\SantaBarbara}{\UCSB} \newcommand{\UCSD}{Center for Astrophysics \& Space Sciences, University of California, San Diego, 9500 Gilman Drive, La Jolla, CA 92093, USA} \newcommand{\IllinoisAstro}{Astronomy Department, University of Illinois at Urbana-Champaign, Urbana, IL 61801, USA } \newcommand{\Cambridge}{Statistical Laboratory, DPMMS, University of Cambridge, Wilberforce Road, Cambridge, CB3 0WB, UK} \newcommand{\KICPCambridge}{Institute of Astronomy and Kavli Institute for Cosmology, Madingley Road, Cambridge, CB3 0HA, UK} \newcommand{\STScI}{Space Telescope Science Institute, 3700 San Martin Dr., Baltimore, MD 21218, USA} \newcommand{\Berkeley}{Department of Astronomy, University of California, Berkeley, CA 94720-3411, USA} \newcommand{\Miller}{Miller Senior Fellow, Miller Institute for Basic Research in Science, University of California, Berkeley, CA 94720, USA} \newcommand{\KICPChicago}{The Kavli Institute for Cosmological Physics, Chicago, IL 60637, USA} \newcommand{\Moore}{The Gordon and Betty Moore Foundation, 1661 Page Mill Road, Palo Alto, CA 94304} \newcommand{\Rutgers}{Department of Physics and Astronomy, Rutgers, the State University of New Jersey, Piscataway, NJ 08854, USA} \newcommand{\UCSC}{Department of Astronomy and Astrophysics, University of California, 1156 High St., Santa Cruz, CA 95064, USA} \correspondingauthor{J.~D.~R.~Pierel} \email{[email protected]} \author{J.~D.~R.~Pierel} \affil{Department of Physics and Astronomy, University of South Carolina, 712 Main St., Columbia, SC 29208, USA} \author{S.~Rodney} \affil{Department of Physics and Astronomy, University of South Carolina, 712 Main St., Columbia, SC 29208, USA} \author{A.~Avelino} \affil{Harvard-Smithsonian Center for Astrophysics, Cambridge, MA 02138, USA} \author{F.~Bianco} \affil{Center for Cosmology and Particle Physics, New York University, New York, NY 10003, USA} \author{A.~V.~Filippenko} \affil{Department of Astronomy, University of California, Berkeley, CA 94720-3411, USA} \affil{Miller Senior Fellow, Miller Institute for Basic Research in Science, University of California, Berkeley, CA 94720, USA} \author{R.~J.~Foley} \affil{Department of Astronomy and Astrophysics, University of California, 1156 High St., Santa Cruz, CA 95064, USA} \author{A.~Friedman} \affil{Center for Astrophysics \& Space Sciences, University of California, San Diego, 9500 Gilman Drive, La Jolla, CA 92093, USA} \author{M.~Hicken} \affil{Harvard-Smithsonian Center for Astrophysics, Cambridge, MA 02138, USA} \author{R.~Hounsell} \affil{Department of Astronomy and Astrophysics, University of California, 1156 High St., Santa Cruz, CA 95064, USA} \affil{Astronomy Department, University of Illinois at Urbana-Champaign, Urbana, IL 61801, USA } \author{S.~W.~Jha} \affil{Department of Physics and Astronomy, Rutgers, the State University of New Jersey, Piscataway, NJ 08854, USA} \author{R.~Kessler} \affil{The Kavli Institute for Cosmological Physics, Chicago, IL 60637, USA} \author{R.~P.~Kirshner} \affil{Harvard-Smithsonian Center for Astrophysics, Cambridge, MA 02138, USA} \affil{The Gordon and Betty Moore Foundation, 1661 Page Mill Road, Palo Alto, CA 94304} \author{K.~Mandel} \affil{Statistical Laboratory, DPMMS, University of Cambridge, Wilberforce Road, Cambridge, CB3 0WB, UK} \affil{Institute of Astronomy and Kavli Institute for Cosmology, Madingley Road, Cambridge, CB3 0HA, UK} \author{G.~Narayan} \affil{Space Telescope Science Institute, 3700 San Martin Dr., Baltimore, MD 21218, USA} \author{D.~Scolnic} \affil{The Kavli Institute for Cosmological Physics, Chicago, IL 60637, USA} \author{L.~Strolger} \affil{Space Telescope Science Institute, 3700 San Martin Dr., Baltimore, MD 21218, USA} \begin{center} \end{center} \begin{abstract} Empirical models of supernova (SN) spectral energy distributions (SEDs) are widely used for SN survey simulations and photometric classifications. The existing library of SED models has excellent optical templates but limited, poorly constrained coverage of ultraviolet (UV) and infrared (IR) wavelengths. However, both regimes are critical for the design and operation of future SN surveys, particularly at IR wavelengths that will be accessible with the James Webb Space Telescope (JWST) and the Wide-Field Infrared Survey Telescope (WFIRST). We create a public repository of improved empirical SED templates using a sampling of Type Ia and core-collapse (CC) photometric light curves to extend the Type Ia parameterized SALT2 model and a set of SN Ib, SN Ic, and SN II SED templates into the UV and near-IR. We apply this new repository of extrapolated SN SED models to examine how future surveys can discriminate between CC and Type Ia SNe at UV and IR wavelengths, and present an open-source software package written in Python, \textit{SNSEDextend}, that enables a user to generate their own extrapolated SEDs. \end{abstract} \ \label{sec:results} The primary output of this work is the open-source \textit{SNSEDextend}\ software package, which is written in Python \citep{Pierel:2018a}\footnote{Code: \href{https://github.com/jpierel14/snsed}{github.com/jpierel14/snsed}; Documentation: \href{http://snsedextend.readthedocs.io}{snsedextend.readthedocs.io} }. The package makes extensive use of the SNCosmo Python package for fitting, calculations, its template library, and more. \textit{SNSEDextend} \ is fully documented, and gives users the capability to generate color curves as in Section \ref{sub:curves} with their own data, and use their own SN templates to produce SED extrapolations into the UV and IR. Using the new open-source \textit{SNSEDextend} \ package, we have produced an initial set of 47 SEDs for SNe~II, Ib, and Ic, plus an extrapolated version of the SALT2 model for SNe~Ia (Figures \ref{fig:allSEDs}, \ref{fig:salt2-extended}). The fully documented and complete repository of these SEDs can be found at an online repository\footnote{\href{dx.doi.org/10.5281/zenodo.1250492}{DOI:10.5281/zenodo.1250492}}. \ \begin{figure} \centering \includegraphics[width=.5\textwidth]{allSEDs} \caption{\label{fig:allSEDs} A representative set of the results produced by the \textit{SNSEDextend} package, for CC~SNe at peak brightness. In each panel the original SED flux density (black) has been normalized, and the shaded regions represent the flux through each bandpass that has been set by the \textit{SNSEDextend}\ package to match the broad-band colors measured in Sections \ref{sub:curves} and \ref{sub:bb}. } \end{figure} \begin{figure} \centering \includegraphics[width=.5\textwidth]{salt2-extended} \caption{\label{fig:salt2-extended} The fully extrapolated SALT2 model, which combines the SALT2-IR model (red) derived in Section \ref{sub:ia_nir} and the SALT2-UV model (blue) derived in Section \ref{sub:ia_uv}. The original SALT2 model (black) was previously defined over the wavelength range 2000--9200~\mbox{\normalfont\AA} \ (though had increased error below 3500~\mbox{\normalfont\AA} \ and above 8000~\mbox{\normalfont\AA} \ [\cite{Guy:2007}]), and now extends over the wavelength range 1700--25,000~\mbox{\normalfont\AA}.} \end{figure} Time-domain science is a key driver in the design of the next generation of observatories, including LSST, JWST, and WFIRST. These telescopes will provide the community with thousands of new SN observations, many of which will include rest-frame UV and NIR photometry. \citep[e.g.,][]{Mesinger:2006,Oguri:2010,najita:2016,Hounsell:2017}. The WFIRST mission, for example, could observe as many as 8,000 SNe at $z<0.8$ with filters covering a rest-frame wavelength range of $\sim7000$--20,000~\mbox{\normalfont\AA}\ \citep{Hounsell:2017}. The new repository of empirically derived SED extrapolations presented here is immediately useful for simulations of these future SN surveys, which are used to optimize survey strategies. The huge number of SN discoveries coming in the next decade will necessitate photometric classifications \citep[e.g.,][]{Kessler:2010,Sako:2011,Campbell:2013,Jones:2018}, as it will not be possible to perform spectroscopic follow-up observations for each object. The extension of the SN SED template library into rest-frame UV and NIR wavelengths may be especially valuable for photometrically isolating SNe~Ia because of their distinguishing spectrophotometric features in those regions. At UV wavelengths, SNe~Ia exhibit a distinct flux deficit relative to CC~SNe \citep{Riess:2004a,Milne:2013}. In NIR bands, SNe~Ia are considered to be excellent standard candles \citep{Dhawan:2017} and exhibit a distinctive secondary maximum in their broadband light curves \citep[e.g.,][]{Leibundgut:1988,Ford:1993}. \begin{figure} \includegraphics[width=.5\textwidth]{classify_redshift_F200WF277W_grid}\caption{\label{fig:jwst} The evolution of CC~SN and SN~Ia colors defined by the JWST F200W and F277W filters (effective wavelengths of $\sim 20,000$~\mbox{\normalfont\AA}\ and $\sim27,700$~\mbox{\normalfont\AA}, respectively), at multiple epochs. Regions where there is separation between the SN~Ia color and all CC~SN colors correspond to redshifts at which distinguishing between SNe~Ia and CC~SNe at that redshift should be possible with the new extrapolated SEDs. The shaded region shows the redshift range over which the nonextrapolated SEDs would have provided color information, and the nonshaded region shows what has been made accessible by the extrapolated SEDs.} \end{figure} A full examination of the accuracy and efficiency of the improvement to photometric classifications is beyond the scope of this work. We have, however, made a preliminary exploration of how the extrapolated SED library will affect color-based classifications with JWST, WFIRST, and LSST (Figures \ref{fig:jwst}, \ref{fig:wfirst}, \ref{fig:lsst}). Using UV-NIR colors at peak brightness (an approach similar to the single-epoch classifications of \citet{Poznanski:2007}), the next generation of telescopes will now be able to distinguish a SN~Ia from a CC~SN with 95\% confidence over 47\% (JWST), 47\% (WFIRST), and 44\% (LSST) of the redshift ranges shown in Figures \ref{fig:jwst}, \ref{fig:wfirst}, and \ref{fig:lsst}, respectively. These rates are a significant improvement upon what was previously possible over the same redshift ranges: 21\% (JWST), 25\% (WFIRST), and 37\% (LSST). To continue improving the extrapolated SED templates, the most valuable additions will be from well-sampled NIR light curves of CC~SNe and UV light curves of SNe~Ia. Aided by the publicly available \textit{SNSEDextend}\ software package, new photometric time-series data can be easily adopted to update the SED templates. These improvements can propagate into more informative simulations for the optimization of future SN surveys, and more accurate photometric classification tools for the analysis of the thousands of SNe those surveys will deliver. \begin{figure} \includegraphics[width=.5\textwidth]{classify_redshift_wfirstjwfirstf_grid} \caption{\label{fig:wfirst} The evolution of CC~SN and SN~Ia colors defined by the WFIRST $J$ and $F$ filters (effective wavelengths of $\sim 12,900$~\mbox{\normalfont\AA}\ and $\sim18,500$~\mbox{\normalfont\AA}, respectively), at multiple epochs. Regions where there is separation between the SN~Ia color and all CC~SN colors correspond to redshifts at which distinguishing between SNe~Ia and CC~SNe at that redshift should be possible with the new extrapolated SEDs. Filter transmission functions were taken from \cite{Hounsell:2017}, which in turn references the WFIRST Cycle 6 instrument parameter release. The shaded region shows the redshift range over which the nonextrapolated SEDs would have provided color information, and the nonshaded region shows what has been made accessible by the extrapolated SEDs. } \end{figure} \begin{figure} \includegraphics[width=.5\textwidth]{classify_redshift_lsstrlssty_grid} \caption{\label{fig:lsst} The evolution of CC~SN and SN~Ia colors defined by the LSST $R$ and $Y$ filters (effective wavelengths of $\sim 6200~\mbox{\normalfont\AA}$ and $\sim9700~\mbox{\normalfont\AA}$, respectively), at multiple epochs. Regions where there is separation between the SN~Ia color and all CC~SN colors correspond to redshifts at which distinguishing between SNe~Ia and CC~SNe at that redshift should be possible with the new extrapolated SEDs. The shaded region shows the redshift range over which the nonextrapolated SEDs would have provided color information, and the nonshaded region shows what has been made accessible by the extrapolated SEDs.} \end{figure} \bigskip \ | 18 | 8 | 1808.02534 |
1808 | 1808.00531_arXiv.txt | The Central Compact Object (CCO) in the Cassiopeia A supernova remnant is most likely a very young ($\approx 300$\,yr) neutron star. If a previously reported decrease of its surface temperature by 4\% in 10 years could be confirmed, it would have profound theoretical implications for neutron star physics. However, the temperature decrease was inferred from \emph{Chandra} ACIS data affected by instrumental effects which could cause time-dependent spectral distortions. Employing a different instrument setup which minimizes spectral distortions, our 2006 and 2012 \emph{Chandra} spectra of the CCO did not show a statistically significant temperature decrease. Here, we present additional observations from 2015 taken in the same instrument mode. During the time span of 8.5 years, we detect no significant temperature decrease, using either carbon or hydrogen atmosphere models in the X-ray spectral fits. Our conservative $3\sigma$ upper limits correspond to $<3.3$\% and $<2.4$\% temperature decrease in 10 years for carbon atmosphere model fits with varying or constant values of the absorbing hydrogen column density, respectively. The recently revised model for the ACIS filter contaminant has a strong effect on the fit results, reducing the significance of the previously reported temperature and flux changes. We expect that a further improved contaminant model and longer time coverage can significantly lower the upper limits in the future. | \label{intro} One of the methods to investigate the composition, structure and physical properties in the interior of neutron stars is to study their thermal evolution (e.g., \citealt{Page2004, Yakovlev2004}). Using \emph{Chandra} observations of the Central Compact Object (CCO) in the Cassiopeia A (Cas A) supernova remnant, and fitting the CCO spectrum with a carbon atmosphere model, \citet[HH10 hereafter]{HeinkeHo2010} reported an unexpectedly rapid 4\% ($5.4\sigma$) decline of the surface temperature and a 21\% flux decline over the time span of 10 years. This rapid cooling was interpreted by \citet{Shternin2011} and \citet{Page2011} as due to enhanced neutrino emission caused by the recent onset of neutron superfluidity (formation of Cooper pairs) in the neutron star core. Considered as the first direct evidence that superfluidity and superconductivity occur in superdense matter of neutron stars, this result has been widely discussed (over 100 publications in 2011-2018). However, the rapid cooling was inferred from \emph{Chandra} ACIS-S Graded mode observations that suffered from several instrumental effects. The most important one is photon pileup, where two or more photons are detected as a single event\footnote{For more details, see \url{cxc.harvard.edu/ciao/ahelp/acis_pileup.html}}. Pileup can distort the observed CCO spectrum. The pileup fraction in the observed spectrum of a given constant source decreases over time because the decreasing sensitivity of the ACIS detector. This is mostly due to an accumulating contaminant on the optical-blocking filters of the ACIS detectors. In addition, not all X-ray events are telemetered in the Graded mode\footnote{For more details, see \url{cxc.harvard.edu/ciao/why/cti.html}}, potentially also affecting the spectrum. For these reasons, the spectral changes reported by HH10 required confirmation. Analyzing observations with different \emph{Chandra} instruments, \citet[E+13 in the following]{Elshamouty2013} reported a statistically significant decrease again only in the case of the ACIS-S Graded mode observations where the best-fit decay was $3.5\%\pm 0.4$\% (from 2000 to 2010). Avoiding the spectral distortion effects caused by photon pileup and the use of the Graded telemetry mode, \citet{PavlovL2009} and \citet{Posselt2013} (P+13 in the following) employed a more suitable instrument mode of ACIS-S in 2006 and 2012 to probe the spectral evolution of the Cas A CCO. Using hydrogen and carbon atmosphere models, P+13 reported that the statistical significance of any temperature change between 2006 and 2012 did not exceed $2.5\sigma$, at the default calibration, for all the considered constraints on the fitting parameters. However, the time coverage was only six years, and the uncertainties were too large to completely rule out the previously reported ``rapid cooling''. Here, we report on new observations with the same ACIS instrument mode, extending the time coverage of the CCO monitoring with this more suitable ACIS observing mode to $8.5$ years.\\ The accumulating ACIS contaminant complicates the analysis of all the CCO \emph{Chandra} ACIS data. Errors in the contamination correction can lead to an offset of the derived spectral parameters from the correct values. P+13 evaluated the influence of the time-variable optical depth of the contaminant on the spectral fit results. An imperfect contamination correction impacts the inferred absorbing hydrogen column density, $N_{\rm H}$, which in turn is correlated with the inferred temperature. Based on the apparently increasing best-fit values of $N_{\rm H}$ from 2006 to 2012, P+13 speculated that an underestimated optical depth of the ACIS contamination could explain any remaining spectral changes of the Cas A CCO. The ACIS Calibration team developed a new model for the ACIS contamination \citep{Plucinsky2016} which is available since December 2016.\footnote{\url{cxc.harvard.edu/caldb/downloads/Release_notes/CALDB\_v4.7.3.html\#TD\_ACIS\_CONTAM\_10}} Here, we also report on the effect of this new contamination model on the results from the 2006 and 2012 subarray data.\\ It is important to emphasize that all the temperature changes were only found if a carbon atmosphere model was used to describe the surface emission of the CCO. Only for this atmosphere model the assumption of a constant emission area is reasonable. This assumption, however, imposes an additional constraint on the spectral fit. When P+13 allowed the emission area to vary, the result was an (insignificantly) \emph{increased} temperature for both carbon and hydrogen atmosphere models. As P+13 showed, hydrogen atmosphere models fit the CCO spectra equally well. The currently existing X-ray data do not allow one to differentiate between these two atmosphere models (see also \citealt{Alford2017}). The carbon atmosphere model was preferred by HH10 and E+13 because it implies an emission size consistent with what one would expect for the entire surface of a neutron star, while the hydrogen atmosphere models results in a much smaller emission area. This could indicate the presence of one or more hot spots. However, one would expect X-ray pulsations in such a case, but none have been found so far (see, e.g., \citealt{PavlovL2009,Mereghetti2002}) though the derived upper limits on the pulsed fraction are above the values measured for other CCOs (e.g., \citealt{Gotthelf2013} and references therein). \begin{figure*}[] {\includegraphics[width=170mm]{fig1.eps}}\\ \caption{Images of the CCO and its vicinity from observing epochs P1 (2006), P2 (2012), P3 (2015), and P4 (2015, three days after P3). North is up, East is to the left, all the images at the same spatial scale. Marked in each image are: the inner source extraction region in black (radius of 4 ACIS pixels $1\farcs{97}$ in P1-P3), the annulus region for the background in yellow lines (inner radius of 5 pixels, outer radius of 10 pixels in P1-P3), and the box regions excluded from the respective background regions in yellow. In P4, the CCO looks elongated because it is located at an off-axis angle of $190\arcsec$ from the optical axis.} \label{ExtractRegions} \end{figure*} \begin{figure}[t] {\includegraphics[height=85mm, angle=90]{fig2.eps}} \caption{Effective areas at the position of the CCO target on the ACIS-S3 chip for different epochs: P1 -- black, P2 -- red, P3 -- blue, P4 (center of the chip) -- cyan. The solid lines correspond to the calibration CALDB 4.7.3, the dashed lines to CALDB 4.5.5.1 which was available for our previous work (P+13). It is clearly seen that the contamination for the old data was underestimated compared to the new contamination model.} \label{effis} \end{figure} | The CCO in Cas\,A does not show statistically significant temperature or flux changes if one employs the spectra obtained in three epochs of the \emph{Chandra} ACIS subarray mode observations which are best suited to such spectral analysis. This conclusion holds for carbon atmosphere as well as hydrogen atmosphere models. An updated calibration, in particular an updated model of the contaminant on the ACIS optical blocking filter, reduced apparent temperature or flux differences for the 2006 and 2012 data to less than $1\sigma$. This confirms the hypothesis by P+13 that an imperfect ACIS contamination model affected the previous findings. While the 2015 subarray data on the CCO imply lower temperature and flux values than in 2012, the overall cooling rates still remain consistent with 0. The current conservative $3\sigma$ \emph{lower} limits on the characteristic cooling times of 300\,yr (varying $N_{\rm H}$) and 410\,yr (tied $N_{\rm H}$) exclude the previously reported rapid cooling by HH10 and E+13 ($\tau_{\rm cool}= 270$\,yr), which was inferred from data obtained using the whole ACIS chip in Graded mode. Because those data of the Cas A CCO are impacted by a substantial pileup effect (distorting the spectra), while the subarray data show negligibe pileup, the theoretical results based on those data (e.g., \citealt{Negreiros2018,Burgio2018,Taranto2016,Grigorian2016,Page2011, Shternin2011}) should be revised, and the new result should be used for any theoretical constraints in neutron star cooling models. \facility{Chandra (ACIS)}\\ \software{CIAO (v4.9; \citealt{Fruscione2006}), XSPEC (v12.8.2; \citealt{Arnaud1996})} | 18 | 8 | 1808.00531 |
1808 | 1808.09626_arXiv.txt | Employing Solar Dynamics Observatory (SDO)/Atmospheric Imaging Assembly (AIA) multi-wavelength images, we report the coronal condensation during the magnetic reconnection (MR) between a system of open and closed coronal loops. Higher-lying magnetically open structures, observed in AIA 171\,\AA~images above the solar limb, move downward and interact with the lower-lying closed loops, resulting in the formation of dips in the former. An X-type structure forms at the interface. The interacting loops reconnect and disappear. Two sets of newly-reconnected loops then form and recede from the MR region. During the MR process, bright emission appears sequentially in the AIA 131\,\AA~and 304\,\AA~channels repeatedly in the dips of higher-lying open structures. This indicates the cooling and condensation process of hotter plasma from $\sim$0.9 MK down to $\sim$0.6 MK, and then to $\sim$0.05 MK, also supported by the light curves of the AIA 171\,\AA, 131\,\AA, and 304\,\AA~channels. The part of higher-lying open structures supporting the condensations participate in the successive MR. The condensations without support by underlying loops then rain back to the solar surface along the newly-reconnected loops. Our results suggest that the MR between coronal loops leads to the condensation of hotter coronal plasma and its downflows. MR thus plays an active role in the mass cycle of coronal plasma because it can initiate the catastrophic cooling and condensation. This underlines that the magnetic and thermal evolution has to be treated together and cannot be separated, even in the case of catastrophic cooling. | \label{sec:int} Magnetic reconnection (MR), the reconfiguration of magnetic field geometry, is a fundamental process in magnetized plasma systems throughout the universe, such as accretion disks, solar and stellar coronae, planetary magnetospheres, and laboratory plasmas \citep{prie00}. It is considered to play an essential role in the rapid release of magnetic energy and its conversion to other forms, e.g. thermal, kinetic, and particles \citep{prie86,prie00}. In solar physics, numerous theoretical concepts of MR have been used to explain various features, e.g. flares, filament eruptions, and coronal mass ejections \citep{shib99,lin00,mei17}. In two-dimensional models, MR occurs at an X-point where anti-parallel magnetic field lines converge and reconnect \citep{prie86,prie00}. The process of MR is difficult to observe directly. However, coronal structures, e.g. loops, and their structural changes often outline the magnetic field topology and its evolution, because the magnetic flux is frozen into the coronal plasma \citep{su13,li16a}. So far, using remote-sensing observations, many MR signatures have been reported, e.g. cusp-shaped post-flare loops \citep{tsun92,yan18}, loop-top hard X-ray sources \citep{masu94,su13}, MR inflows \citep{yoko01,li09,yang15,huang18} and outflows \citep{taka12,tian14,li16c,ning16}, supra-arcades downflows \citep{mcke00, inne03, sava12, li16b}, current sheets \citep{liu10,kwon16,li16a,li16b,li18,xue18}, and plasmoid ejections \citep{kuma13,yang17,yang18,zheng17}. The condensation of cool plasma out of the hot corona is a widely observed phenomenon best seen at the solar limb. One widely investigated concept for this process is based on the thermal properties of the plasma alone \citep{mull03}. It is independent of the (evolution of the) coronal magnetic field, and only the loss of equilibrium between heat input, heat conduction and radiative losses causes the plasma to cool catastrophically \citep{xia16}. In numerical models, thermal instabilities occurring within a current sheet and an accompanying MR result in the formation of a quiescent (current sheet) prominence \citep{smith77,malh83}. On the other hand, MR can create a helical flux rope, and the cooling radiation and condensation of plasma trapped inside the flux rope form the cool dense plasma of prominences \citep{pneu83,kane15}. Recently, \citet{kane17} propose a MR-condensation prominence formation model, and demonstrate that MR can lead to flux rope formation and radiative condensation under certain conditions. Employing Solar Dynamics Observatory \citep[SDO,][]{pesn12}/Atmospheric Imaging Assembly \citep[AIA,][]{lemen12} images, the cooling and condensation of coronal plasma are observed in a loop system \citep{liu12} and a coronal cavity \citep{berg12} for prominence formations, and in an active region loop \citep{vash15} for coronal rain formation, respectively. In this Letter, we study the MR between coronal loops, and firstly report the cooling and condensation of coronal plasma during the MR. The observations and results are presented in Sections\,\ref{sec:obs} and \ref{sec:res}, and a summary and discussion is presented in Section\,\ref{sec:sum}. | \label{sec:sum} Employing SDO/AIA multi-wavelength images, we study the evolution of the MR between two sets of coronal loops L1 and L2, and the following condensation of hotter coronal plasma in the dip region of loops L1. The loops L1 show curved open structures with a dip in AIA 171\,\AA~images. They move downward and reconnect with the lower-lying closed loops L2. Newly-reconnected loops L3 and L4 form, and retract away from the MR region. Along with the MR, bright emission appears sequentially in the AIA 131\,\AA~and 304\,\AA~channels. The cooling and condensation of hotter coronal plasma takes place repeatedly in the dip of loops L1. The time delays between the peaks of the light curves of the AIA 171\,\AA~and 131\,\AA~channels, and the AIA 131\,\AA~and 304\,\AA~channels, are 30 and 50\,min, respectively. Due to the successive MR, the condensation, without support by lower-lying loops L1, flows toward the chromosphere first along the left leg, and then along the right leg, of the newly-reconnected loops L4. According to the AIA extreme-ultraviolet (EUV) observations and the PFSS coronal magnetic field, schematic diagrams are provided in Figure\,\ref{f:cartoon} to describe the evolution of MR and coronal condensation. A set of open field lines (L1) come close to a region with closed field lines (L2), see Figure\,\ref{f:cartoon}(a), forming a dip in the former. Between them MR takes place, and newly-reconnected field lines (L3 and L4) form, see Figure\,\ref{f:cartoon}(b). Coronal material is gathered in the dip and cools, see Figure\,\ref{f:cartoon}(b). Through a loss of thermal equilibrium triggered by the MR, a condensation forms, and then slides down along the newly-formed field lines (L4), see Figure\,\ref{f:cartoon}(c). The convergence of loops L1 and L2 toward an X-structure, and the retraction of newly-formed loops L3 and L4 away from the X-structure show clear observational evidence of MR occurring in the X-structure. The MR lasts for $\sim$22\,hr, much longer than any other MR events abundantly encountered on the Sun, which typically last only a few minutes, such as explosive events, EUV bursts, and jets, to a few hours, such as in flares. Both loops L1 and L2 are rooted in quiet Sun regions, with weaker magnetic field. Because much less magnetic flux reconnects in a much longer time, one can imagine that the MR rate here is much smaller, consistent with the observations that no associated flaring activity is detected during the MR in our case. Large coronal condensations are rarely observed. In this study, the coronal plasma cools from the AIA 171\,\AA~channel, rather than other AIA higher temperature channels, e.g. 193\,\AA~ \citep{liu12} and 211\,\AA~\citep{berg12}. The time delay between the AIA 171\,\AA~and 304\,\AA~channels is 80\,min, similar to \citet{vash15}, but less than \citet{liu12} and \citet{berg12}. Moreover, the AIA 131\,\AA~channel is employed to better show the cooling process of coronal plasma. Condensations of plasma along open field lines are found. So far coronal rain has been reported to occur only along preexisting structures that are magnetically closed \citep{anto01,schr01}. When considering the AIA 304\,\AA~images alone, the condensations here would also resemble that occurring in closed loops. However, the source region for the condensation is the magnetically open structure L1 that forms a dip where the condensation takes place. In the later stages the condensation rains down not only along the leg of that open structure L1, but also through the MR region and down along the leg of the newly-reconnected field lines L4. This provides a new and alternative mechanism for the formation of coronal rain away from the magnetically closed regions that is initiated by MR. In traditional models, the thermal evolution, i.e. condensation, and the evolution of the magnetic field are treated separately \citep{mull03,xia16}. Mostly, it is assumed that the condensations form through a loss of equilibrium between heating and cooling. If the field line is stretched in the horizontal direction, and the condensation contains sufficient mass, then the field line would form a dip and plasma will remain in the lower density corona within the dip \citep{karp01}. Alternatively, MR can cause a helical structure with numerous dips during the prominence formation \citep{pneu83,kane15}. Considering the energy balance along these structures, one then gets condensations and plasma is trapped, again, in these dips \citep{pneu83,kane15,kane17}. In this Letter, both dips and condensations in the dips are formed, and they would resemble the prominence formation models \citep{pneu83,kane15}. However, the condensation remains for only a short time, and then rains down to the chromosphere when the dips are broken by the successive MR. In our study, we show that the MR and the condensation cannot be treated separately, but that plasma condensation naturally arises during the MR process. \begin{figure}[ht!] \centering \plotone{f1.eps} \caption{Evolution of the coronal emission and magnetic field lines from a PFSS model. ((a)-(c)) AIA 171 \AA~images. ((d)-(e)) Coronal magnetic field lines derived from the PFSS model. The red rectangle on the white circle in (a) indicates the location of FOV of ((a)-(c)) with respect to the full Sun. The AIA images are rotated counter-clockwise by an angle of 40$^{\circ}$. The white and red arrows N and W in (a), (b), and (d) show the north and west directions in the FOVs of AIA before and after the rotation, respectively. The cyan rectangle in (a) shows the FOV of Figure\,\ref{f:condensation}, and the blue rectangle in (b) marks the region for the light curve of the AIA 171\,\AA~channel as displayed in Figure\,\ref{f:signatures}(a) by the blue line. The green, blue, red, and cyan dotted lines in (a) and (c) outline the coronal loops L1, L2, L3, and L4, and the red solid arrows in (a) and (d) mark a dip of loops L1. The red and blue lines AB and CD in (b) and (c) separately indicate the positions of time-slices of AIA 171\,\AA~images as displayed in Figures\,\ref{f:signatures}(a) and (b). In (d) and (e), the blue and green lines represent the closed and open PFSS coronal magnetic field lines at 00:04:00 UT (d) and 06:04:00 UT (e), respectively, and the inner gray-scale images show the line of sight (LOS) magnetograms in a range of $\pm$100 Mx cm$^{-2}$. In (e), the outer image displays the LASCO/C2 coronagraph image at 09:48:05 UT, and the dashed line describes the upper radial boundary of the PFSS model corona, also called the source surface, at 2.5 solar radii. \label{f:reconnection}} \end{figure} \begin{figure}[ht!] \plotone{f2.eps} \centering \caption{Temporal evolution of the coronal emission near the MR region. Time-slices of AIA 171\,\AA~images along the red and blue lines AB (a) and CD (b) in Figures\,\ref{f:reconnection}(b) and (c), respectively. In (a), the blue, green, and red lines separately show the light curves of the AIA 171\,\AA, 131\,\AA, and 304\,\AA~channels in the blue and yellow rectangles in Figures\,\ref{f:reconnection}(b), \ref{f:condensation}(e), and \ref{f:condensation}(i). The cyan vertical lines denote the times of Figures\,\ref{f:reconnection}((a)-(c)). The green arrows mark the fine structures of loops L1. The cyan dashed line outlines the downward motion of loops L1. In (b), the green dotted lines show the contractions of loops L3, and the cyan dotted lines outline the contractions of loops L4. The blue dotted line denotes the backward motion of loops L4. The moving speeds of loops L1, L3, and L4 are denoted by the numbers in ((a)-(b)). \label{f:signatures}} \end{figure} \begin{figure}[ht!] \plotone{f3.eps} \centering \caption{Evolution of the condensation from $\sim$1\,MK to below 0.1\,MK. AIA 171\,\AA~((a)-(c)), 131\,\AA~((d)-(f)), and 304\,\AA~((g)-(i)) images. The yellow rectangles in (e) and (i) mark the regions for the light curves of the AIA 131\,\AA~and 304\,\AA~channels as displayed in Figure\,\ref{f:signatures}(a) by the green and red lines, respectively. The red and green ellipses in (b), (e), and (h) enclose the condensation of coronal plasma in (h). \label{f:condensation}} \end{figure} \begin{figure}[ht!] \centering \includegraphics[width=0.5\textwidth]{f4.eps} \caption{Temporal evolution of the condensation. ((a)-(b)) AIA 304\,\AA~images. ((c)-(d)) Time-slices of AIA 304\,\AA~images separately along the blue lines EF (c) and GH (d) in (a) and (b). The green dotted lines in (c) and (d) outline the downflows of the condensations. \label{f:downflow}} \end{figure} \begin{figure}[ht!] \centering \includegraphics[width=\textwidth]{f5.eps} \caption{Spatial relation between condensation downflows and loops. AIA 171\,\AA~(a), 131\,\AA~(b), and 304\,\AA~(c) images, and their composite (d), with the same FOV as in Figures\,\ref{f:reconnection}((a)-(c)) and \ref{f:downflow}((a)-(b)). (An animation of this figure is available.) \label{f:animation}} \end{figure} \begin{figure}[ht!] \centering \includegraphics[width=0.5\textwidth]{f6.eps} \caption{Schematic diagrams of the MR and coronal condensation. In ((a)-(c)), the grey thick lines denote the solar limb. The green, blue, and green-blue lines separately show the magnetic field lines of loops L1, L2, L3, and L4, whose directions are marked by the red arrows. In (a), the green arrow denotes the dip of loops L1 and also the moving direction. In (b) and (c), the purple stars show the MR between loops L1 and L2. The green, yellow, and red patches indicate the AIA 171\,\AA, 131\,\AA, and 304\,\AA~plasma, respectively. \label{f:cartoon}} \end{figure} | 18 | 8 | 1808.09626 |
1808 | 1808.10051_arXiv.txt | {} {In this work we present an algorithm to identify compact groups (CGs) that closely follows Hickson's original aim and that improves the completeness of the samples of compact groups obtained from redshift surveys.} {Instead of identifying CGs in projection first and then checking a velocity concordance criterion, we identify them directly in redshift space using Hickson-like criteria. The methodology was tested on a mock lightcone of galaxies built from the outputs of a recent semi-analytic model of galaxy formation run on top of the Millennium Simulation I after scaling to represent the first-year Planck cosmology.} { The new algorithm identifies nearly twice as many CGs, no longer missing CGs that failed the isolation criterion because of velocity outliers lying in the isolation annulus. The new CG sample picks up lower surface brightness groups, which are both looser and with fainter brightest galaxies, missed by the classic method. A new catalogue of compact groups from the Sloan Digital Sky Survey is the natural corollary of this study. The publicly available sample comprises $462$ observational groups with four or more galaxy members, of which $406$ clearly fulfil all the compact group requirements: compactness, isolation, and velocity concordance of all of their members. The remaining $56$ groups need further redshift information of potentially contaminating sources. This constitutes the largest sample of groups that strictly satisfy all the Hickson's criteria in a survey with available spectroscopic information.} {} | Over the past 40 years, the astronomical community has devoted much time to the systematic search for compact groups (CGs). These peculiar small galaxy systems have proven to be a very powerful tool to understand galaxy interactions in dense environments, shaping our knowledge of galaxy evolution. The most popular sample of CGs was identified by \cite{Hickson82}. This sample was a systematic search of CGs on plates of the Palomar Observatory Sky Survey that relied on a visual inspection. \citeauthor{Hickson82} established a set of rules that should be fulfilled for a galaxy system to be considered a compact group: compactness, isolation (with a relatively empty annulus surrounding the group), and population. These rules were defined to obtain small, isolated, and compact galaxy systems. In this way, \citeauthor{Hickson82} created the well-known Hickson Compact Group (HCG) sample that comprises 100 galaxy groups identified in projection on the sky. However, since only angular positions were used to identify the sample, their truly compact nature was questioned (e.g. \citealt{Mamon86,Mamon90}). A few years later, the redshifts of their galaxy members were measured and a velocity filter was added to the criteria, resulting in a sample of 68 CGs with at least four concordant velocities \citep{Hickson92}. Since then, several works have identified CGs in observational or simulated galaxy catalogues based on Hickson's recipes with or without including the velocity filtering depending on the available information (e.g. \citealt{Prandoni+94,Iovino+02,Lee+04,McConnachie+08,McConnachie+09,DiazGimenez&Mamon10,DiazGimenez+12,sohn+15}). Most of these works share the same philosophy for the algorithm construction, i.e. finding CGs that meet Hickson's criteria through an automatic procedure that resembles the original Hickson's visual inspection by first detecting candidate CGs on the sky, and then discarding obvious foreground/background galaxies along the line of sight from their discordant redshifts. However, for quasi-complete spectroscopic galaxy catalogues this procedure is not optimal because it may discard groups that appear not isolated on the sky, although the galaxies populating the isolation annulus turn out to all have discordant redshifts. Having chance projected galaxies in the cone is not a problem if subsequent velocity filtering is performed. The big problem is that if such galaxies lie outside the CG, they will lead the algorithm to discard the group as non-isolated, whereas they are clear chance projections along the line of sight. The type of restrictions that arises from the procedure itself could easily bias subsequent studies related to CGs environment. \begin{figure*} \begin{center} \includegraphics[width=9.1cm,trim=125 90 90 120,clip]{fig_1a} \includegraphics[width=9.1cm,trim=130 155 120 120,clip]{fig_1b} \caption{\label{flux_diagram} Flow charts showing the automatic implementation of the Hickson's criteria to identify CGs. On the left the flow chart shows the \texttt{classic} algorithm, while on the right the \texttt{modified} algorithm is shown. } \end{center} \end{figure*} In the past, some authors already noticed this problem and suggested new ways to identify CGs in redshift surveys by changing the searching procedure and using percolation algorithms similar to the friends-of-friends (FoF) algorithm (e.g. \citealt{Barton96,Focardi&Kelm02,Zandivarez03,sohn+16}). However, most of these attempts disregarded most or all of Hickson's definitions, and only kept a compactness criterion based on the physical size of the galaxy systems or the intergalactic separations. The purpose of the present work is to present a \texttt{modified} algorithm that applies the Hickson's criteria directly in redshift space. As a result, we present a new catalogue of CGs extracted from the Sloan Digital Sky Survey (SDSS) that satisfies Hickson's criteria. The layout of this work is as follows. In Section 2 we describe the construction of the mock galaxy lightcone used in this work for the testing process. This catalogue was constructed using synthetic galaxies obtained from a semi-analytical model (SAM) of galaxy formation \citep{Henriques+15}. In Section 3 we describe the well-known Hickson's recipe for identifying CGs \citep{Hickson82,Hickson92} and the previous and new algorithms for identifying CGs that fulfil these conditions. We apply these algorithms to the mock catalogue and perform the corresponding comparisons between the resulting CG catalogues. In Section 4 we apply both algorithms to an observational galaxy sample compiled from the SDSS by \cite{tempel17} and compare the resulting new CG sample with previous observational identifications. Finally, in Section 5 we summarise our results and present our conclusions. | \label{sec:discus} The original definition of compact groups states that they are small, high density associations of bright galaxies that are relatively isolated in space \citep{Hickson82}. To identify such systems, groups must meet several criteria: limited population within a magnitude range, compactness, spatial isolation within a magnitude range, and velocity concordance of all of their galaxy members. While in principle the limiting values that can be adopted for each of the criteria are arbitrary, it is customary to adopt the definitions introduced by \cite{Hickson82} and \cite{Hickson92}. In this work, we followed these original ideas, and adopted the commonly used limits to present a new algorithm to identify CGs in redshift surveys. The algorithm has been tested using a mock lightcone built from galaxies extracted from a recent semi-analytical model of galaxy formation \citep{Henriques+15} run on top of the Millennium Simulation I rescaled to represent the cosmological model determined by the Plank cosmology. Given the methodology used to test the algorithm, the choice of a different semi-analytical model will not affect the results presented in this work. To test the new algorithm we compared the resulting sample with that obtained using a previous version of the algorithm that had been used in the past to create and analyse samples of CGs \citep{DiazGimenez&Mamon10,DiazGimenez+12,DiazGimenez+15,Taverna+16}. We called the previous application of the algorithm \texttt{classic}, since it basically reproduces the steps followed by \citeauthor{Hickson82}: it starts looking for CG in projection in the plane of the sky, and only when the projected sample is available does it check whether all the galaxy members have concordant radial velocities. Other authors have also followed these steps \citep{MendesdeOliveira&Hickson91,McConnachie+09,sohn+15}. The main disadvantage of the \texttt{classic} path is that when the algorithms involve the counting of galaxies in a magnitude range within a region of the projected sky, it may include many interlopers that lie in the line of sight in the background or foreground of the group itself. These interlopers will affect the population, as well as the isolation of the potential groups, and eventually some of the groups may be excluded from the projected sample. This could thus compromise the completeness of the resulting sample. With the issue of the completeness in mind, we introduced a \texttt{modified} version of the algorithm. Instead of applying the velocity filter on a pre-selected projected sample of CGs, the velocity concordance is a requirement imposed as the galaxies are being included as galaxy members or neighbours. As a result, we obtained modified samples that are roughly twice the size of the classic samples. We have shown that, compared to the \texttt{classic} version, the \texttt{modified} searching of CGs includes more groups towards the limit of the compactness criterion, with larger projected sizes, with fainter first-ranked galaxies, and with more similar two first-ranked galaxies. The incompleteness of the classic sample shows no dependence on distances to the groups. Therefore, our \texttt{modified} algorithm helped us to improve the completeness of the samples of CGs. An interesting question now arises about the purity of the samples. Using the information in 3D real space from the simulation, and following \cite{DiazGimenez&Mamon10}, we investigated the occurrence of chance alignments in our classic and modified CG samples. We found that % the algorithm does not affect the fraction of CGs that are chance alignments of galaxies along the line of sight. It would have been desirable to diminish the fraction of chance alignment groups, but at least it has not increased even though we doubled the number of identified CGs. In a future work, we will investigate how to increase the fraction of physically dense compact groups (i.e. diminish the fraction of chance alignment groups) via observational constraints. As a corollary of this study, we present a new CG catalogue identified on the observational spectroscopic galaxy sample compiled by \cite{tempel17} from the SDSS DR12. This CG catalogue comprises $462$ systems. After performing visual and automatic inspections of the surroundings of each CG in the photometric SDSS sample, we determined that $406$ CGs ($\sim 88\%$ of the sample) can certainly be considered isolated, while for the remaining $56$ CGs we were not able to classify them as surely isolated, due to possible photometric contamination in their surrounding area caused by incompleteness inherent to the SDSS spectroscopic catalogue (saturation of bright galaxies and fibre collision). Further spectroscopic information is needed. In addition, we performed a detailed comparison with the available sample of CGs identified by \cite{sohn+16}. The comparison between samples of CGs identified by different authors, algorithms, and criteria is never straightforward, and one has to be careful when extracting general conclusions about compact groups that could stand for CGs defined in one way but do not hold for a different sample. It is important to have larger samples of CGs identified in a unique homogeneous way, with well known selection criteria to obtain statistically reliable conclusions. With this aim either the sample of \cite{sohn+16} and the sample introduced in this work achieve this goal, and to shed light on the properties of small peculiar galaxy systems, although the two samples differ in their definitions of what the authors understand by compact groups. Particularly, criteria regarding flux limit, isolation and embedded compact groups are key in determining the properties of the final sample of CGs. The sample of 462 Hickson Modified Compact Groups (\texttt{HMCG}s) presented in this work has become the largest catalogue of CGs that satisfy all of Hickson's original criteria: they are small, compact, and isolated associations of four or more concordant galaxies. This sample is available for the astronomical community as electronic table in this publication (see Appendix~\ref{catalogue} for details). | 18 | 8 | 1808.10051 |
1808 | 1808.09283_arXiv.txt | Studying the stellar mass, age, luminosity, star-formation rate, and impact parameter of quasar absorber host galaxies can aid in the understanding of galaxy formation and evolution as well as in testing their models. We derive the Spectral Energy Distribution (SED) and impact parameter limits of low redshift ($z_{abs} = 0.37 - 0.55$) Mg II absorbers and of higher redshift ($z_{abs} = 1.0 - 2.5$) 2175 \AA\ dust absorbers (2DAs). We use an imaging stacking technique to statistically boost the signal-to-noise ratio (SNR) to increase detection of the absorber host galaxies. The point spread function of the background quasar is modeled with Principal Component Analysis (PCA). This method efficiently reduces the uncertainty of traditional PSF modeling. Our SED for Mg II absorbers indicates that low redshift Mg II absorber host galaxies are likely star-forming galaxies transitioning into red quiescent galaxies, with a low star formation rate of 2.2 $M_\odot$ $yr^{-1}$. From the stacked images and simulations, we show that the average impact parameter of 2DAs is > 5 times smaller than that of Mg II absorbers, at < 7 kpc instead of Mg II absorbers' 48 kpc, indicating that 2DAs are likely associated with disk components of high redshift galaxies. This means that 2DAs are likely good probes to study precursors to the Milky Way. | Galactic evolution is a cornerstone of astronomy that is being widely and intensely studied all over the world. One of the most powerful ways to gain insight into the properties of high-redshift galaxies is by tracking metal absorption lines in interstellar clouds. Quasars, powered by the accretion of materials in supermassive black holes, are the brightest objects in the early universe. They are able to probe metal absorption lines in intervening galaxies between the quasars and Earth, making quasar absorption line systems one of the most powerful tools in studying the gas content in the early universe to high levels of sensitivity that do not depend on redshift. Bahcall \& Spitzer (1969) suggested that metal absorption lines seen in quasar spectra are induced by large gas halos surrounding host galaxies. These gas halos can extend up to 100 kpc away. Thus, tracking quasar absorption lines to study these gas halos has been an active field for the past few decades. Previous studies that investigated host galaxies of quasar absorbers have used deep imaging while the gas content has been extensively studied with spectroscopic follow-ups. However, such studies are extremely restrictive, due to small sample sizes and high expenses in performing spectroscopic follow-ups of faint objects. Thus, most studies in the late 1990s and early 2000s were limited to a few dozen case studies and focused on individual attributes of varying host galaxies. Furthermore, the wide distribution of redshifts and magnitudes of the sample result in the inability to determine strong statistical properties of the absorbers and their host galaxies. Overarching trends across different ranges of redshifts and absorption strengths could not be found. These trends are necessary for categorizing the major components found in all absorbers, such as average luminosity-weighted impact parameters for investigating location and geometry of absorbers, and average spectral energy distribution for studying stellar population, mass, and star formation history. On the other hand, ground-based telescope imaging is widely available but significantly noisier, making detection of absorber host galaxies in individual frames impossible. One method of boosting the signal-to-noise ratio (SNR), and hence sensitivity, of ground-based imaging is by stacking many frames of imaging data together, relying on statistical properties of large datasets to preserve the signals of each individual frame while decreasing the overall noise level. By the central limit theorem of statistics, the noise distribution in the mean image is tighter than the original images by a factor of $\sqrt{N}$, where N is the number of frames stacked. Stacking images brings out signals that were previously undetectable in individual frames since the background noise level is greatly reduced. This not only mitigates the challenges associated with expensive, deep field, high contrast imaging follow-ups, but also offers a statistical method to study host galaxy properties as a galaxy population. Correlations between different attributes, such as surface brightness, galaxy age, mass and star formation history, can then be determined. Stacking approaches and studies have provided valuable results in a number of areas. For example, Bartelmann \& White (2003) demonstrated that stacking of ROSATA ll-Sky Survey X-ray images of high-redshift clusters detected in the Sloan Digital Sky Survey (SDSS) can be used to derive their mean X-ray properties. In the context of galaxies, Hogg et al. (1997) constrained the IR signal from faint galaxies using stacked Keck data. Similarly, Brandt et al. (2001) measured the mean X-ray flux of Lyman break galaxies, Zibetti et al. (2004) characterized the very low surface brightness ``diffuse" light in galaxy halos, and White et al. (2007) constrained the radio properties of SDSS quasars down to the nanojansky level. In spectroscopic studies, stacking techniques have been extensively used to look for weak signals. Composite spectra of SDSS quasars were used to detect weak absorption lines, as well as dust reddening effects that are well below the noise level in individual spectra (Nestor et al. 2003; M\'{e}nard et al. 2005; York et al. 2006). Clearly, the stacking approach is very advantageous in investigating faint sources (Zibetti et al. 2007). Zibetti et al. (2005a, b, 2007) used such a stacking method for studying large samples of Mg II absorption quasars. The Mg II absorption feature was chosen as their interest of study because of dominant ions in H gas. Mg II possesses a resonance transition at the doublet (2796.35, 2803.53\AA). Zibetti et al. (2007) extracted thousands of Mg II absorber quasar cutouts from the Sloan Digital Sky Survey's Data Release 4 (DR4). They binned the Mg II absorbers by the redshift of the absorbers and by the rest equivalent width (REW) of the 2796 \AA\ absorption line. They then centered the absorbers to their calculated centroids by way of interpolation in order to make the images super-imposable, subtracted the quasar's point spread function (PSF) to diminish the quasar's light contributions, and finally rescaled the intensity of the absorbers to achieve normality and consistency. Zibetti et al. (2007) corrected the images for galactic extinction and mean-stacked the images to reveal one final mean frame. Their results indicate that there is no significant redshift dependence for both impact parameter and rest-frame colors for redshifts up to $z_{abs} = 1$. They also showed that stronger absorption systems display the colors of blue star-forming galaxies while weaker absorption systems mostly originate from red passive galaxies. Finally, Zibetti et al. (2007) demonstrated the stacking technique's usefulness in detecting the light of QSO hosts and their environments. Since 2007, there have been numerous advances in cataloging quasar absorber host galaxies. Bouche et al. (2007) discovered that there may exist a correlation between star formation rate and the equivalent width of the Mg II absorption doublet, indicating that gas content can possibly be used as a predictor for starburst phenomena. Chen et al. (2010) used 70 low impact parameter Mg II absorbers from the Magellan Echellette spectrograph to determine that increased impact parameters of the absorbers result in decreased absorber strength. Kacprzak et al. (2012) used 88 spectroscopically confirmed Mg II absorbers from the Hubble Space Telescope and SDSS to demonstrate that there exists an azimuthal bimodal distribution of absorbers, with blue star-forming galaxies driving the bimodality. They indicated that halo gas exists more commonly around the projected galaxy's major and minor axis. In addition, the bimodality is generated by the accretion of gas along the galaxy's major axis and outflowed along the galaxy's minor axis. This is consistent with galaxy evolution scenarios where star formation galaxies accrete gas. In these same scenarios, red galaxies typically have smaller star formation rates due to lower gas reservoirs. Nielsen et al. (2013) cataloged 182 spectroscopically identified, high redshift ($0.7 - 1.1$) Mg II absorbers, indicating that low stellar mass galaxies tend to be bluer and hence possess higher star formation rates. However, they also indicated that the Mg II absorption is preferentially weaker in such systems. As shown, all of these studies required a lot of dedicated spectroscopic data acquired from telescopes and used a small sample size. A recently discovered type of absorber that is of great interest is the quasar 2175 \AA\ dust absorbers (2DAs, Wang et al. 2004, Ma et al. 2017, 2018a). These dust absorbers, displaying strong broad 2175 \AA\ absorption features and dust extinction (e.g., Wang et al. 2004; Zhou et al. 2010; Jiang et al. 2010a, b, Jiang et al. 2011; Wang et al. 2012; Zhang et al. 2015; Pan et al. 2017; Ma et al. 2015; 2017), closely resemble the Milky Way (MW). The similarity between the extinction curve of these absorbers and that of the MW means that studying these systems might provide clues on the evolution of the MW. 2DAs are also a subgroup of Mg II absorbers, representing about 1 per cent of the total Mg II absorber population detected (Zhao et al. 2018, in prep). While extensive studies on Mg II absorbers have been conducted, as shown above, very few studies on 2DAs and their host galaxies have been made. Ma et al. (2017, 2018a) performed correlation analysis between metallicity, velocity width, redshift, depletion level, and other quantities to compare 2DAs with other absorption lines. They concluded that the 2DA host galaxies contain high metallicity, high depletion, are generally massive, and are chemically enriched. 2DAs are also more massive than typical Damped-Lyman-alpha (DLA) and sub-DLA galaxies, showing greater maturation and age. The median estimated stellar mass of 2DA host galaxies is $2 \times 10^{10}$ $M_\odot$ (Ma et al. 2018a). Thus, studying 2DAs will also give insight into the dust attenuation in high redshift galaxies, which is consequential in the theory of galactic evolution. Early results (Ma et al. 2017, 2018a, b) show that 2DAs are likely matured and massive. They also report that the impact parameter for one 2DA system at $z_{abs} = 2.12$ is only 5.5 kpc, or 0.65 arcsec away from the quasar. However, this is only an isolated case as acquiring HST time is extremely difficult and it is hard to study a large number of absorber host galaxies using HST. For a sample of 436 2DAs identified by us in the SDSS-III DR 12 imaging data (Zhao et al. 2018, in prep), HST imaging for each individual absorber system is cost inhibitive. Therefore, only stacking is a viable option that will result in the sensitivity necessary for detection and for studying the overarching statistical properties of all 2DA systems. Our research is to develop an automatic and fast imaging stacking and subtraction method using data from ground-based telescopes to study the statistical properties of host galaxies of Mg II absorbers and 2DAs i.e. their average impact parameter and SEDs. The impact parameter will possibly reveal the geometry of host galaxies and will also help identify which components of host galaxies are observed through quasar absorption. In order to study the host galaxy properties and impact parameter distributions of the newly discovered and extremely faint 2DAs, we create a stacking technique that is specialized in reducing noise levels as effectively as possible. We use better and more accurate stellar profile approximation methods to decrease the noise level and boost detection. We validate our technique by comparing our results with Zibetti et al.'s (2007). Because of the large datasets, we create an automated package for ease of processing. This package can also be used to perform stacking analysis on other types of absorbers, such as Ca II, CIV or Fe II absorption systems. Our goal is to identify aspects of galaxy evolution, discover any new properties/correlations, and confirm previous studies' results by probing the statistical properties of Mg II absorbers and 2DAs. Because the 2DAs are largely precluded in the imaging data due to the high redshift range ($z_{abs} = 1.0 - 2.5$), our improvements to noise-reduction in the stacking procedure will aid in revealing properties of 2DAs. Hopefully, this will result in new findings about the MW in the future. In Section 2 of this paper, we present our methodology of collecting imaging data and selecting both Mg II absorbers, 2DAs, and their corresponding reference QSOs. In Section 3 we present the image processing techniques in terms of subtracting the sky background, masking unwanted sources, calibrating image intensities, etc. In Section 4 we calculate surface brightness (SB) profiles and classify the derived SED of Mg II absorber host galaxies. We also perform simulations of the 2DA host galaxies' average impact parameter. In Section 5 we discuss our findings. In Section 6 we compare our technique to previous studies. | In this study, we have shown that PCA is capable of reproducing quasar profiles to a high degree of accuracy. We now explore the differences in methodology. Zibetti et al. (2007) utilized a masking algorithm that covered up all sources deemed unlikely to be an absorbing galaxy at the given redshift. They used SExtractor as their baseline masking, supplemented by a flux-limited mask derived from a metal-poor stellar population. We used the fpObj files released by DR7 and DR12, which contain detailed categorizations of the sources found in each field image, ensuring that our masking masks all stars, extended sources (such as field galaxies, nebula), etc. Next, during the PCA fitting of the quasars, we chose to only take the components calculated in a $13 \times 13$ square around the centroid of the quasar. Although the fitting is highly accurate in most cases, some reference quasars leave significant amounts of residue after masking and PSF subtraction. The $13 \times 13$ was chosen to minimize overfitting based on the average extension of a non-absorbing quasar's immediate residues, which are calculated from their FWHM. Although the procedure was highly effective in retaining data (> 425 out of 436 passed in all bands, with failed cases being where the QSO was unable to be cut due to boundaries or no suitable sources were found for PCA PSF fitting), the amount of residue left over is low; the surrounding noise level is still very high. Two conclusions were made after comparison between the two absorber stackings: \begin{enumerate} \item The high redshift of the 2DAs ($z_{abs} = 1 - 2.5$) means that the detection completion is very low. Assuming a standard cosmological model, magnitudes in that range increase by approximately $2 - 4$ from the magnitudes recorded at $z_{abs} = 0.37 - 0.55$ due to cosmological dimming. Thus, the SNR is already much weaker due to increased distance. Flux-limited masks that helped distinguish between Mg II absorber host galaxies and non-absorbing galaxies were not necessary for 2DAs. \item The impact parameters for most of the 2DA host galaxies are small: < 0.8 arcsec or < 7 kpc from the quasar. Since the light of the QSO in SDSS imaging extends out to a few kpc, the 2DAs are almost completely precluded in the imaging data. \end{enumerate} | 18 | 8 | 1808.09283 |
1808 | 1808.05252_arXiv.txt | We present a neutral hydrogen-selected absorption-line survey of gas with \hi\ column densities $15 < \mlnhi < 19$ at $z\la 1$ using the Cosmic Origins Spectrograph on the {\em Hubble Space Telescope}. Our main aim is to determine the metallicity distribution of these absorbers. Our sample consists of \ssz\ absorbers selected on the basis of their \hi\ absorption strength. Here we discuss the properties of our survey and the immediate empirical results. We find singly and doubly ionized metal species and \hi\ typically have similar velocity profiles, implying they probe gas in the same or similar environments. The ionic ratios (e.g., \ncii/\nciii, \noi/\ncii) indicate the gas in these absorbers is largely ionized, and the ionization conditions are quite comparable across the sampled \nhi\ range. The Doppler parameters of the \hi\ imply $T \la 5\times 10^4$ K on average, consistent with the gas being photoionized. The \mgii\ column densities span $>2$ orders of magnitude at any given \nhi, indicating a wide range of metallicities (from solar to $<1/100$ solar). In the range $16.2\la \mlnhi \la 17$, there is a gap in the \nmgii\ distribution corresponding to gas with $\sim 10\%$ solar metallicity, consistent with the gap seen in the previously identified bimodal metallicity distribution in this column density regime. Less than $3\%$ of the absorbers in our sample show no detectable metal absorption, implying truly-pristine gas at $z\la 1$ is uncommon. We find $\langle [$\feii/\mgii$]\rangle = -0.4 \pm 0.3$, and since $\alpha$-enhancement can affect this ratio, dust depletion is extremely mild. | Modern theory and cosmological simulations agree that the star formation of galaxies and the properties of their circumgalactic medium (CGM) should be intimately connected \citep[see recent review by][]{tumlinson17}. This is especially true for the flows through the CGM: feedback from star formation is understood to drive outflows that carry mass and metals away from galaxies, while infall from the intergalactic medium (IGM) is thought to bring in fresh gas to fuel on-going star formation. Without significant feedback, most baryons would cool into the centers of halos to form prodigious quantities of stars \citep[e.g.,][]{white78,keres09}, but with feedback, the baryon content of stars and cold gas in galaxies can be matched by driving matter into the CGM and beyond \citep[e.g.,][]{fukugita98,conroy09}. Similarly, without continued infall of IGM material, star-forming galaxies would consume their interstellar gas in $\sim$1 Gyr (e.g., \citealt{rocha-pinto00}). Gas accretion may also control, in part, the evolution of the elemental abundances in galaxies and could also play a role in the mass-metallicity relationship \citep[e.g.,][]{kacprzak16}. These exchanges of matter through the CGM thus play critical roles in the evolution of galaxies. In our Cosmic Origins Spectrograph (COS) CGM Compendium (CCC) presented in this paper, we aim to directly explore how a specific property of the CGM --- the metallicity --- is distributed throughout cosmic time and environments. Observationally, outflows have been characterized at various redshifts and appear ubiquitous in the universe \citep[e.g.,][]{shapley03,steidel10,weiner09,martin13,rubin14}. However, direct observational evidence for cold gas accretion (colder than the virial temperature, i.e. gas at $ \sim 10^{4}$--$10^{5}$ K) has been more difficult to discover. There is some evidence for gas accretion onto the Milky Way and other nearby galaxies (e.g., \citealt{wakker01,fraternali08,fox16,fox18}). Recent integral-field spectroscopic observations and studies of redshift-space distortions also suggest some evidence of inflows at low and high redshifts \citep[e.g.,][]{fumagalli16,bielby17,cantalupo17,turner17}. The $z=0$ observations indicate that gas can be accreted via other means than cold flow accretion, while the higher redshift results provide some support to cold flow accretion as observed in simulations, although none of these observational results show yet conclusively and unambiguously the existence of IGM gaseous filaments feeding galaxies. Another indirect way to probe cold flow accretion (and more generally gas flows in and out of galaxies) is via QSO absorbers that probe the CGM of galaxies. According to cosmological simulations \citep[e.g.,][]{fumagalli11a,vandevoort12a,vandevoort12b,faucher-giguere15,hafen17}, the covering fraction of the cold streams in these simulations appears to peak in a \hi\ column density (\nhi) regime known as the (partial) Lyman Limit systems (pLLSs and LLSs). Using QSO spectra, we can search specifically for these absorbers and empirically characterize their properties. In this and subsequent papers, we adhere to the following definition for the various absorbers studied in our survey. The pLLSs and LLSs have \hi\ column densities $16.2 \la \mlnhi <17.2 $ and $17.2 \le \mlnhi <19$, respectively. The $\mlnhi = 17.2$ cutoff corresponds to an optical depth at the Lyman limit $\tau_{\rm LL} = 1$ ($\tau_{\rm LL} \equiv \mnhi\ \sigma_{\rm H\,I}$, where $\sigma_{\rm H\,I} = 6.30 \times 10^{-18}$ cm$^2$ is the absorption cross section of a hydrogen atom at the Lyman limit, see \citealt{spitzer78}). The 16.2 dex lower bound of the pLLSs is more arbitrary, but corresponds to $\tau_{\rm LL} = 0.1$, which creates a break at the Lyman limit that is still visible in moderate and high signal-to-noise ratio (SNR) spectra (see \citealt{lehner13}, hereafter \citetalias{lehner13}). This definition of LLSs differs from our previous surveys but follows more closely the generally adopted definition of LLSs and is motivated by the finding of \citet{wotta16} (hereafter \citetalias{wotta16}), which shows a tentative difference in the metallicity distribution between the pLLSs and LLSs at $z\la 1$ (see below). We also use the standard definition for the damped \lya\ absorbers (DLAs) that have $\mlnhi \ge 20.3$ and super-LLSs (SLLSs, a.k.a. sub-DLAs) with \nhi\ in the range $19.0 \le \mlnhi <20.3$. The absorbers with $15 \la \mlnhi <16.2$ have no proper definition and we simply define these absorbers as the strong \lya\ forest systems (SLFSs). We, however, emphasize the SLFSs are more related to the diffuse CGM than the IGM at $z<1$. Indeed, several studies of the galaxy--absorber two-point cross-correlation function have shown significant clustering between galaxies and SLFSs while a weak or absent clustering signal for weaker \nhi\ absorbers \citep[e.g.,][]{lanzetta95,penton02,bowen02,chen05,prochaska11a,prochaska11c,tejos14}, pointing to a strong physical connection between SLFSs (and other stronger \hi\ absorbers) and galaxies. While the diffuse gas probed by SLFSs, pLLSs, LLSs cannot be yet directly imaged, the properties of these absorbers can inform us on their origins and hence help us characterize the gas in the CGM of galaxies at low redshift. One of the key properties is the metallicity since the enrichment levels of the gas help in differentiating between the plausible origins of the gas (\citealt{ribaudo11,fumagalli11a}; \citetalias{lehner13}). For example, we can differentiate accretion from the IGM (that is metal-poor) from the more metal-rich gas produced by galaxies (which we emphasize can be outflowing or inflowing). Determining the metallicity distribution of the CGM gas also informs us about the fraction that has remained mostly untouched by the successive episodes of star formation in galaxies over billions of years. While metallicity estimates nearly always require large ionization corrections in the $15 < \mlnhi < 19$ range, several independent studies have shown that such corrections can be well constrained by the broad range of ions despite an uncertain EUV ionizing background (\citealt{howk09}; \citetalias{lehner13}; \citealt{crighton13a,fumagalli16}; \citetalias{wotta16}). The uncertainty in the ionizing background can lead to an absolute uncertainty in the metallicity of a factor 2--3 (\citealt{howk09}; \citetalias{wotta16}; \citealt{fumagalli16}). Though it is not precision cosmology, it is accurate enough to separate low from high metallicities as well as enable a determination of the metallicity probability distribution function of these absorbers. We bear in mind that this error level on the metallicity of the absorbers is comparable to the magnitude of uncertainty on chemical abundances determined from emission lines in galaxy spectra (e.g., \citealt{berg16}). However, we also note that the uncertainties in the metallicities for a given ionizing background are also typically quite small (depending on the exact constraints provided by the available ions; see \citetalias{lehner13}; \citealt{fumagalli16}). Over the last few years, our team has led several surveys of the metallicities of the pLLSs and LLSs at $z\la 1$ (\citealt{lehner09}; \citetalias{lehner13}; \citealt{ribaudo11,tripp11}; \citetalias{wotta16}) and $z\sim 2$--4 (\citealt{lehner16}, and see also \citealt{fumagalli16,cooper15,glidden16}). One of our main findings has been the first evidence of widespread low metallicity gas in the pLLS and LLS regime that is not observed in higher column density absorbers at similar redshifts. At $z\la 1$ there is indeed a clear prevalence of very low metallicity pLLSs with $\xh \equiv \log N_{\rm X}/N_{\rm H} - \log {\rm X/H}_\sun \la -1.4$ (where X is an $\alpha$-element such as Mg, O, Si) with $43\% \pm 8\%$ of the pLLSs with these low metallicities, while $<10\%$ of the DLAs and SLLSs have $\xh < -1.4$ (where X is now a $\alpha$-element, Zn, or Fe corrected for depletion or $\alpha/{\rm Fe}$ enhancement) \citepalias{wotta16}. The pLLSs and LLSs therefore uniquely probe metal-poor gas in relatively dense (denser than the IGM probed by weak \lya\ forest absorbers) environments in the universe at low and high redshift (\citetalias{lehner13,wotta16}; \citealt{lehner16,fumagalli16,cooper15,glidden16}). Another major finding from these initial surveys was that the shape of the metallicity distribution of the pLLSs at $z\la 1$ appears to be bimodal with an equal number of pLLSs below and above $\xh = -1$, a functional form only observed at $z<1$ (at $z>2$, the metallicity distribution appears to be unimodal for similar absorbers, see \citealt{lehner16}). These empirical results provide the first strong evidence of widespread very low metallicity gas around $z\la 1$ galaxies that may possibly accrete onto galaxies, although we emphasize that there is not yet any direct evidence that this gas is actually accreting. Our initial survey consisted of 28 absorbers (23 pLLSs and 5 LLSs) \citepalias{lehner13}. Our second survey doubled the size of our initial sample (44 pLLSs and 11 LLSs), confirming our initial results, and tentatively demonstrating a possible change in the shape of the metallicity distribution between the pLLSs and LLSs and a lower frequency of metal poor LLSs compared to pLLSs \citepalias{wotta16}. Our surveys have increased by an order of magnitude the number of pLLSs and LLSs where the metallicities have been estimated at $z\la 1$ compared to the status prior to the installation of the COS onboard the {\it Hubble Space Telescope} (\hst) (see compilation by \citealt{lehner09} and references therein). However, the sample of absorbers is still relatively small, especially when we treat the pLLSs and LLSs separately. Furthermore to effectively probe the level of the IGM/CGM enrichment requires that we not only increase the number of strong \hi\ absorbers, but that we also target lower \nhi\ absorbers. We were awarded a {\it Hubble Space Telescope} (\hst) Legacy program to produce the largest survey to date of \hi-selected absorbers at $z\la 1$ using the spectra observed with the high-resolution mode (G130M and G160M gratings) of COS available at the \textit{Barbara A. Mikulski Archive for Space Telescopes} (MAST). This archive is the richest database of UV QSO spectra with sufficient quality (spectral resolution and SNR) that can permit accurate estimations of the column densities of metal and \hi\ transitions. As we detail in this paper, \nhi\ can be well determined by modeling the entire Lyman series (which is accessible for $z_{\rm em}\ga 0.2$ QSOs in the COS bandpass). The restframe UV, FUV, and EUV wavelengths covered by the COS spectra include a large number of metal lines with a wide range of strengths and ionization states, and additional NUV transitions (\mgii\ $\lambda \lambda$2796, 2803, \feii\ $\lambda \lambda$2382, 2600) can be observed from ground-based telescopes for the redshifts probed by our survey. At $z\la 1$, we note that any absorbers with $\mnhi \ga 10^{15.2}$ \cmm\ are in a relative ``sweet spot'' for an unbiased metallicity study because both the metal and \hi\ column densities can be accurately estimated to provide a reliable census of the metallicity distribution. The SNRs of the COS spectra are, however, not high enough to pursue an unbiased survey of the metallicity in the more diffuse gas probed by \lya\ forest absorbers with $\mlnhi < 15$ at $z\la 1$. Here we present our CCC survey of \hi-selected absorbers with column densities in the range $15<\mlnhi < 19$ at $z\la 1$ aimed to determine the metallicities of these absorbers. We used the \hst\ COS G130M and G160M archive that we supplemented with additional spectra of \mgii\ and \feii\ obtained from the High Resolution Echelle Spectrometer (HIRES) on Keck I and the Ultraviolet and Visual Echelle Spectrograph (UVES) on the Very Large Telescope (VLT). Our COS G130M and G160M survey presented here comprises \ssz\ absorbers, and our total sample for the metallicity study has \tsz\ absorbers (the additional absorbers that were primarily observed with other instruments or gratings). In this paper, we present the design survey, observations, and data reductions (\S\ref{s-present}), the search for strong \hi\ absorption in the COS spectra (\S\ref{s-search}) and column density and kinematics measurements of the \ssz\ absorbers (\S\ref{s-estimate-col}), and the immediate implications from these measurements (\S\ref{s-results}). We summarize our main results in \S\ref{s-summary}. In the subsequent papers, we will present the grids of ionization models combined with a Bayesian formalism and Markov Chain Monte Carlo (MCMC) techniques to robustly determine the metallicities (and errors) of the absorbers (paper II), the metallicity probability distribution of the pLLSs and LLSs (paper II) and SLFSs (paper III), and the properties of the high ions in these absorbers (paper IV). | \label{s-summary} Using the \hst\ COS G130M and G160M archive, we have built the largest sample to date of \hi-selected absorbers with $15<\mlnhi < 19$ at $z\la 1$ for which we can estimate the metallicities. The sample analyzed in this paper has \ssz\ absorbers, and our total sample consists of \tsz\ absorbers (where the additional data were observed \hst\ STIS, \fuse, and \hst\ COS). This survey is about an order of magnitude larger than the first survey of pLLSs and LLSs at $z<1$ we undertook \citepalias{lehner13}, and it is the first survey targeting and analyzing absorbers with $15<\mlnhi < 16.2$ in the same manner as the pLLSs and LLSs. More quantitatively, we increase the sample of pLLSs by a factor 2 (from 44 in \citetalias{wotta16} to 82) and of LLSs by a factor 2.6 (from 11 in \citetalias{wotta16} to 29), and assemble for the first time a sample of 152 SLFSs. The \hi\ selection ensures that no bias is introduced in the metallicity distribution of these absorbers. We consider only absorbers with $\mlnhi >15$ because the spectra of $z<1$ QSO do not have high enough SNRs to sensitively probe metallicity that is $\la 10\%$ solar at lower column densities. For about half of our sample observed with COS G130M and/or G160M (\gbmgii\ absorbers), we have also high resolution Keck HIRES and VLT UVES observations of the strong NUV transition of \mgii\ and \feii. All the measurements and spectra are made available online. In subsequent papers, we will present the photoionization models used to determine the metallicity of these absorbers, the metallicity distributions of these absorbers, and the evolution of the metallicities over 7 orders of magnitude in \nhi\ ($15 < \mlnhi \la 22$). We will also study how the properties of the high ions (in particular \ovi) correlate with those of the low ions (paper IV). Here we have presented the basic measurements (column densities and kinematics) from the absorption of these absorbers to empirically characterize some of the properties of gas with $15<\mlnhi < 19$ at $z\la 1$. Our initial main findings are as follows. \begin{enumerate}[wide, labelwidth=!, labelindent=0pt] \item From the comparison of the absorption profiles and the width of the profiles, we conclude low ions (singly and doubly ionized species) and \hi\ trace the same gas. Thus they can be modeled {\it a priori} with a single ionization phase model. This does not necessarily apply to the higher ions, and absorbers with high ion absorption must be multiphase. From the profile fitting of the \hi\ transitions, we show that on average the gas temperature is cool with $T < 5 \times 10^4$ K, which is consistent with the gas being photoionized. \item The [\mgii/\hi] ratio spans over 2 orders of magnitude at any \nhi\ over the range $15< \mlnhi < 19$, implying a metal enrichment from a solar to $\la$1/100 solar metallicity. We find that the absorbers are most likely all metal-enriched at some level, with only $<$3\% of the absorbers (90\% confidence level) showing no metal absorption in the spectra. The bimodal metallicity distribution observed for the pLLSs in \citetalias{lehner13} and \citetalias{wotta16} is apparent from the \mgii\ column densities and equivalent widths, with a lack of data in the range $16.2\la \mlnhi \la 17.2$ around a metallicity corresponding to about 10\% solar. This corresponds to the previously observed dip in the metallicity distribution. \item There is no strong dependence of the ionic ratios with \nhi. This implies that the ionization properties of the absorbers must not change dramatically over the range $15<\mlnhi < 19$ (4 orders of magnitude in \nhi). \item We estimate $\langle [$\feii/\mgii$]\rangle = -0.4 \pm 0.3$ comparable to that observed in DLAs at any $z$. This ratio can be affected by $\alpha$-enhancement from Type II supernovae, which implies that Fe dust depletion is small if any. There is also no trend of the [\feii/\mgii] ratio with \nhi, which might be expected if dust depletion was a dominant factor. These findings demonstrate that depletion onto dust in these absorbers is negligible. On average we also show that [C/$\alpha$] is consistent with a solar value, but with a large scatter of about 0.3 dex. \end{enumerate} | 18 | 8 | 1808.05252 |
1808 | 1808.06067_arXiv.txt | Quasi-periodic oscillations (QPO) originating from innermost regions of blazars can provide unique perspective of some of the burning issues in blazar studies including disk-jet connection, launch of relativistic jets from the central engine, and other extreme conditions near the fast rotating supermassive black holes. However, a number of hurdles associated with searching QPOs in blazars e.g., red-noise dominance, modest significance of the detection and periodic modulation lasting for only a couple of cycles, make it difficult to estimate the true significance of the detection. In this work, we report a $\sim$ 330-day QPO in the Fermi/LAT observations of the blazar Mrk 501 spanning nearly a decade. To establish consistency of the result, we adopted multiple approaches to the time series analysis and employed four widely known methods. Among these, Lomb-Scargle periodogram and weighted wavelet z-transform represent frequency domain based methods whereas epoch folding and z-transformed discrete auto-correlation function are time-domain based analysis. Power spectrum response method was followed to properly account for the red-noise, largely inherent in blazar light curves. Both local and global significance of the signal were found to be above 99\% over possible spurious detection. In the context where not many $\gamma$-ray QPOs have been reported to last more than 5 cycles, this might be one of the few instances where we witness a sub-year timescale $\gamma$-ray QPO persisting nearly 7 cycles. A number of possible scenarios linked with binary supermassive black hole, relativistic jets, and accretion disks can be invoked to explain the transient QPO. | Blazars, the most powerful sources in the Universe, are a small sub-class of active galactic nuclei (AGN). They are believed to harbor monstrously giant black holes, ($\sim 10^8-10^{10}M_\odot$), squeezed within a small volume comparable to that of the Solar System \citep[see][for recent black hole imaging of M87 galaxy]{2019ApJ...875L...6E}. The surrounding accretion disk constantly feeds the black hole with tremendous amount of matter which makes the nuclei extremely bright and consequently outshine the whole galaxy. The central engine spews out matter nearly at the speed of the light in the form of relativistic jets directed towards the Earth. Blazars consists two kinds of sources: flat-spectrum radio quasars (FSRQ), the ones which show emission lines and have the synchrotron peak in the lower electromagnetic frequency (typically $< 10^{14}$ Hz); and BL Lacertae (BL Lac) objects, the others which show weak or no emission lines and have synchrotron peak in the higher frequency. Although less powerful than FSRQs, BL Lacs are considered an extreme class of sources with highest synchrotron and inverse-Compton energies. More recently, blazar TXS 0506+056 was associated with the first non-stellar neutrino emission detected by the IceCube experiment \citep{2018Sci...361.1378I,2018Sci...361..147I, Padovani2018}. Blazars being the most luminous sources are \emph{visible} across the universe; and therefore, they can be used to probe the large scale geometry of the universe. Moreover, blazars could form an ideal sample of sources to study some of the challenging issues related to AGN e.g., interaction between accretion disk and jets, and the conditions leading to the launch of jets around supermassive black holes. Although blazar cores mostly at large distances are not resolved by any current instruments, the presence of the relativistic jets in blazars ``highly exaggerates'' the variability amplitudes and the timescales through relativistic beaming effects \citep[see][]{up95} such that information from the central engine can be carried along to us. For the reason, variability studies can be one of the most powerful methods -- if not the only method -- that guide us to the innermost regions of AGN, otherwise completely hidden from our view. \begin{figure*} \begin{center} \includegraphics[width=0.98\textwidth,angle=0]{Fig1.pdf} \caption{The longterm Fermi/LAT (0.1--300 GeV) weekly binned observations of the blazar Mrk 501. The vertical magenta lines approximately mark the peaks of the periodic oscillations (see Section \ref{sec:analysis}).} \label{Fig:1} \end{center} \end{figure*} Although, the flux variability shown by AGN, in general, appears to be aperiodic, the statistical nature of the observed variability can be broadly characterized as \emph{red-noise}, meaning source flux changes by larger amplitude over longer period of time. However, over the past decade presence of quasi-periodic oscillations (QPO) in the multi-frequency blazar light curves has been recorded. The reported periodic timescales range from a few hours to several years \citep[see][and references therein]{Bhatta2017}. In particular, detection of QPOs in the $\gamma$-ray light curves of a small number of blazars are intriguing. Of them, the famous case of 2.18-year periodicity in the TeV blazar PG 1553+113 was first reported in \citealt{Ackermann2015} and more recently in \citet{Tavani2018}. In addition, several works have claimed presence of $\gamma$-ray QPOs in other sources as well: \citet{Sandrinelli14} claimed $\sim1.73$-year QPO in the $\gamma$-ray light curve of the blazar PKS 2155-304 which later was confirmed by \citet{Zhang2017c}. Similarly, \citet{Zhang2017a} reported 3.35-year periodicity in the blazar PKS 0426-380 in 8-year long observation period. \citet{Zhang2017b} claimed 2.1-year QPO in the 9-year long observations of the source PKS 0301-243. In 8-years long $\gamma$-ray observations, \citet{Sandrinelli2017} studied QPO in three blazars S5 0716+714, Mrk 421, BL Lac and concluded that the first two blazars do not show any periodicity while BL Lac exhibited a $\sim$1.86-year QPO, similar to the one found in optical observations the source. It is important to note here that in all of these cases the claimed periodic timescales were nearly two years and above (or low temporal frequency), and number of observed periodic cycles were less than 5. More recently, \citet{Covino2018} studying several blazars in $\gamma$-ray reported no significant periodicities in the light curves. Also \citet{Sandrinelli2016b} studied a number of blazars in the $\gamma$-ray band, including PKS 2155-304, OJ 287, and 3C 279 PKS 1510-089 and claimed the presence of year-like QPOs with modest significance. Similarly, \citet{Prokhorov2017} in their search for the $\gamma$-ray sources showing periodicities confirmed previously reported quasi-periodicities in the three blazars, PG 1553+113, PKS 2155-304 and BL Lacertae, and further observed evidence for quasi-periodic behavior of 4 blazars: S5 0716+716 and three of the high redshift blazars -- 4C +01.28, PKS 0805-07 and PKS 2052-47. Mrk 501 (R.A.=$\rm 16^{h}53^{m}52.2^{s}$, Dec.= $+39\degr45\arcmin37\arcsec$) is one of the nearest blazars (z = 0.034; Ulrich et al. 1975) that shines bright in the X-ray. Together with Mrk 421, it is one of the earliest extragalactic sources detected in the TeV band \citep{Quinn1996}. Based on the location of the synchrotron peak, it is classified as high synchrotron peaked BL Lacs (HSP; \citealt{Abdo2010}). The source has been extensively studied in a broad range of the electromagnetic spectrum by using several instruments including MAGIC, VERITAS, Whipple 10m Telescope, Fermi/LAT, RXTE and Swift, and radio and optical wavelengths \citep[][and references therein]{Ahnen2017,Furniss2015}. A 23-day period was claimed to be found in the multi-frequency observations \citep[see][and references therein]{Rieger2000}, and later a harmonic with 72-day period was also reported by \citet{Rodig2009}. We organize the contents of the paper as following: In Section \ref{sec:obs}, Fermi/LAT data acquisition and analysis are discussed. In Section \ref{sec:analysis}, we present the time series analysis of the blazar Mrk 501 using decade long $\gamma$-ray observations. Although we analyzed the light curve using four methods, to avoid lengthy discussion on methods and put emphasis on more popular methods, we discuss Lomb-Scargle periodogram (LSP) and the weighted wavelet $z$-transform (WWZ) methods in the main paper while we move the epoch folding and z-transformed discrete auto-correlation function analysis to Appendix. In the section, we report a $\sim$ 330-day quasi-periodic modulations in the $\gamma$-ray light curve. We also note that the periodic signal gradually weakened near the end of the observation period. Finally, in Section \ref{discussion} we discuss some of the possible scenarios that can explain the observed periodicity. | \label{discussion} We analyzed $\sim$10 year long Fermi/LAT observations of the famous blazar Mrk 501 using LSP, WWZ, epoch-folding and ZACF methods. All four methods consistently detected a transient quasi-periodic oscillations of the period centered around 332 days. The period in the source rest frame ($P$) is estimated as $P=P_{obs}/(1+z)$ $=321$ days, where $P_{obs}$ and $z$ are the observed period and the red-shift, respectively. The observed transient feature might have arisen due to correlation of the noise e.g. colored (power-law) noise in the light curve. Therefore, we employed both frequency and time domain based analysis, and subsequently, estimated the local and global significance of the observed period taking account a number of artefacts that potentially could produce spurious peaks e.g. red-noise, discrete data, finite observation length, and uneven data sampling. We simulated a large number of light curves from the best-fit PSD model and subsequently assigned the sampling of the real observation. The significance (both local and global) of the detection in LSP and WWZ methods was found to be above 99\%. In other words, the probability that observed periodicity was generated due to correlation of noise was $<1\%$ and therefore, the feature at the timescale of 332 days is highly likely to be of the physical origin. A quick analysis of the source was also performed in the optical waveband using the publicly available data from the Catalina Real-time Transient Survey\footnote{\url{http://nesssi.cacr.caltech.edu/DataRelease/}}. Some hints of the $\gamma$-ray detected periodicity was also found in the analysis. However, the optical observations being sparse with large gaps in between, we could not be conclusive about the results. A multi-frequency time series analysis of the source is deferred to the future study. It should be pointed out that most of the $\gamma$-ray QPOs mentioned in the Introduction are reported to have lasted only a few cycles (typically less 5 cycles). Against such a backdrop, what we have observed might be the first interesting event where we witnessed transient $\gamma$-ray periodic modulations in the high energy emission of the blazar that persist up to 7 cycles before the oscillations gradually fade away. It could be also possible that the flux modulations started long before the Fermi/LAT launch year (i.e. 2008). In that case, the modulations might have spanned more than 7 cycles. In any case, the detection adds Mrk 501 to the list of a few reported cases of QPOs in blazar $\gamma$-ray light curves. Detection of such periodic signals is immensely useful because the signals, mostly originating from the innermost regions of the blazar cores, carry information about the inner structures including space-time geometry around fast spinning supermassive black holes, disk-jet connection, and magnetic field configuration of the accretion disk. QPOs in blazars can be linked to a host of possible events that are chiefly based on three scenarios: supermassive binary black holes (SMBBH), accretion disk based instabilities, and events related to geometric configuration of the jet. Below we discuss some of these in the context of the observed QPO in Mrk 501. \begin{itemize} \item Supermassive binary black holes: Periodic $\gamma$-ray flux modulations are most likely to originate from the relativistic jets, but we may still need a periodic perturbation close to the central engine which would propagate along the jet. It is found that SMBBH systems provide the most natural model for such recurring perturbations \citep[e.g.,][]{Lehto96,Valtonen08} that give rise to periodic flux modulations. In Mrk 501, a 23-day periodic behavior previously reported was based on this model \citep[see][]{Rieger2000}. Interestingly, in both the cases the maxima on average are nearly 8 times the minima of the modulations. The SMBBH orbit can fairly be assumed to be circular as dynamical friction gradually tends to erase any initial eccentricity over the course of merging. The radius of the Keplerian orbit can be calculated using the relation, \begin{equation} r=9.5\times 10^{-5}\left ( \frac{M}{10^9M_\odot } \right )^{1/3}\tau _k^{2/3} \rm \ pc \end{equation} In the case of Mrk 501 which has a black hole of mass $7\times10^8M_\odot$ \citep{Ghisellini2010}, the total mass of the system can safely be assumed to be $1\times10^9 M_\odot $. Then the black holes can be estimated to have a separation of $\sim 5$ milli-parsecs. However, such a close SMBBH systems has not been detected yet. If the jet axis is not aligned to the total angular momentum vector of the binary system, it will induce precession in the trajectory leading to accretion disk precession under gravitational torque \citep[e.g][]{Katz1997} or the jet precession \citep[e.g.][]{Begelman1980,Caproni2017,Sobacchi2017}, and also precession in periodic impact as in OJ 287 \citep{Sillanpaa88}. In the case of significantly elliptical binary orbit, the secondary BH periodically perturbs the primary jet at periastron to induce magnetohydrodynamic (MHD) instabilities. These, in turn, can lead to particle acceleration via magnetic reconnection events. Consequently, the particles emit via synchrotron and inverse Compton emission \citep[see][for details]{Tavani2018}. In the case of Doppler-boosted flux modulations, the true period in the source rest frame takes the form $P_{obs}=(1+z)(1-\beta cos\theta)\tau_k$, so that for a typical bulk Lorentz factor of 10, the corresponding Keplerian orbit at the radius of 0.1 pc has a period of $\sim 90$ years. \citep[For illustration, refer to the left panel of Figure 4 in][]{Bhatta2018d} \item Accretion disk based events: QPOs might arise simply due to hot-spots on the accretion disk revolving around the black holes. In such cases, taking the black hole mass $7\times10^8M_\odot$ for the source, the radius of the corresponding hot-spot turns around to be about 90 gravitational radii $r_g$ or about 4 milli-parsecs, where $r_g=GM/c^2$. Similarly, tilted inner disks around spinning black holes undergo Lense-Thirring precession resulting in periodic flux modulations\citep[e.g.][]{Stella98}. The Lense-Thirring effect around a spinning black hole with Kerr spin parameter $a$ causes tilted disk orbits to precess about its angular momentum vector with a period conveniently expressed as \begin{equation} \tau_{LT}= \frac{1}{2a} \left ( \frac{r}{r_g} \right )^{3} \tau_{k-r_g}, \end{equation} \noindent where $\tau_{k-r_g}$ is the Keplerian period at the black hole's gravitational radius, $r_g = GM/c^2$. In such scenario, precession around fast spinning black hole ($a=0.9$), with the given mass of the source, the observed period in the source rest frame can be used to estimate the inner radius of the orbit to be $\simeq 13\ r_g$. When the disk is strongly coupled to the jet, the precession of the disk can make the jets precess. Recently \citet[][]{Liska2018} performed 3D general relativistic MHD simulations based on the similar scenario; the results showed that the jets exhibit precession with a period of $\sim 1000 t_g$, where $t_g=r_g/c$, which is of the order of the period we observed in the source. Similarly, in the case of globally perturbed thick accretion disks \citep[e.g.][]{Rezzolla2003,Zanotti2003}, the disk undergoes p-mode oscillations such that the fundamental frequency of the oscillation can be approximated as, \begin{equation} f_{0}\approx 100\left ( \frac{r}{r_g} \right )^{-3/2}\left ( \frac{M}{10^8M\odot } \right )^{-1} \ \rm day^{-1} \end{equation} \citep[see][]{Liu2006,An2013}. If the observed 321 rest-frame period is taken to be the period corresponding to such fundamental frequency, the inner radii of the perturbed disks turns out to be $\sim 276\ r_g$. Interestingly, this represents the order of the inner radius of the truncated disk in some of the radio loud bread-line region galaxies showing disk-jet coupling \citep[see e.g.,][]{Bhatta2018b}. In the sources showing strong disk-jet connection, such perturbations could propagate along the jet resulting in high energy QPOs. \item Jet Geometry: The transient 332\--day period might as well be linked to geometrical configuration of the jet, such that instabilities propagating with relativistic speeds along the magnetized jets e.g., following a helical path \citep[as in][]{Camenzind92,Mohan15}, could provide a possible explanation for the QPOs in blazars. In such a scenario, for the emission regions moving with a typical bulk Lorentz factor $\Gamma \sim 10$ along the trajectory viewed about an angle, $\theta \sim 1/\Gamma$, the time period in the source rest frame, corresponding to 332 days observed period, can be estimated using $P\simeq \Gamma ^2P_{obs}/(1+z)$. The calculated timescale is in the order of one 100 years, and it provides us an upper limit of a few tens of parsecs to the spatial scale for the periodic structures. Alternatively, perturbations close to the central engine propagate along the jet inducing periodic swings in the jet viewing angle. Using Doppler boosting relation $F(\nu)=\delta^{(3+\alpha)}F'(\nu')$, for a given change in the angle $\Delta \theta$, the ratio of the observed flux modulation, intrinsic flux remaining unchanged, can be given as \begin{equation} \Delta logF=-\left ( 3+\alpha \right )\delta \Gamma \beta sin\theta \Delta \theta , \end{equation} where negative sign indicates that for a positive change in the angle, representing the emission region moving away from the line of sight, the flux ratio decreases. In such events, periodic changes in the viewing angle as small as 1$^{\circ}$ can also give rise to periodic flux modulations by an order of magnitude \citep[For illustration for flux doubling case, refer to the right panel of Figure 4 in][]{Bhatta2018d}. \end{itemize} We also note the interesting feature in the light curve (see Figure \ref{Fig:1}) that the QPO fades around the same time when the $\gamma$-ray emission diminishes in intensity. This might indicate a close link between QPOs and jet production mechanisms in blazars. Such possible link can still be explained in the scenarios discussed above. For example, the gravitational perturbation in a SMBBH system can feed to the instabilities \emph{in situ} that contribute to the $\gamma$-ray production and consequently, result in transient QPOs; the disk hot-spots influencing jet emission can frequently form and disrupt on some characteristic timescale. Similarly, the jet can turn away reducing the $\gamma$-ray emission and alter the periodicity at the same time via Doppler boosting. Finally, although all of these scenarios can lead to QPOs, without further multi-frequency time series analysis and collective discussion on the topic, the exact cause of the observed transient $\gamma$-ray rhythmic flux modulations would still remain mysterious. | 18 | 8 | 1808.06067 |
1808 | 1808.08550_arXiv.txt | The flux transport dynamo, in which the poloidal magnetic field is generated by the Babcock--Leighton mechanism and the meridional circulation plays a crucial role, has emerged as an attractive model for the solar cycle. Based on theoretical calculations done with this model, we argue that the fluctuations in the Babcock--Leighton mechanism and the fluctuations in the meridional circulation are the most likely causes of the irregularities of the solar cycle. With our increased theoretical understanding of how these irregularities arise, it can be possible to predict a future solar cycle by feeding the appropriate observational data in a theoretical dynamo model. | % The flux transport dynamo model, which started being developed about a quarter century ago \citep{WSN91,CSD95,Durney95}, has emerged as an attractive theoretical model for the solar cycle. There are several modern reviews \citep{Chou11,Chou14, Charbonneau14, Karakreview14} surveying the current status of the field. The present paper is not a comprehensive review, but is based on a talk in a Workshop at the International Space Science Institute (ISSI) highlighting the works done by the author and his coworkers. Readers are assumed to be familiar with the phenomenology of the solar cycle and the basic concepts of MHD (such as flux freezing and magnetic buoyancy). Readers not having this background are advised to look at the earlier reviews by the author \citep{Chou11,Chou14}, which were written for wider readership. The initial effort in this field of flux transport dynamo was directed towards developing periodic models to explain the various periodic aspects of the solar cycle. Once sufficiently sophisticated periodic models were available, the next question was whether these theoretical models can be used to understand how the irregularities of the solar cycle arise. There is also a related question: if we understand what causes the irregularities of the cycle, then will that enable us to predict future cycles? We discuss the basic periodic model of the flux transport dynamo in the next Section. Then \S~3 discusses the possible causes of the irregularities of the solar cycle. Afterwards in \S~4 we address the question whether we are now in a position to predict future cycles. Finally, in \S~5 we summarize the limitations of the 2D kinematic dynamo models and the recent efforts of going beyond such simple models. | We have pointed out that over the years we have acquired an understanding of how the irregularities of the solar cycle arise and that this understanding helps us in predicting future solar activity. Our point of view is that the fluctuations in the Babcock--Leighton mechanism and the fluctuations in the meridional circulation are the two primary sources of irregularities in the solar cycles. These fluctuations have to be modelled realistically and fed into a theoretical dynamo model to generate predictions. It may be noted that we now have a huge amount of data on the magnetic activity of solar-like stars \citep{Chou17}. Some solar-like stars display grand minima and we have evidence for the Waldmeier effect in some of them---see the concluding paragraph of \citet{KKC14}. This suggests that dynamos similar to the solar dynamo may be operational in the interiors of solar-type stars and the irregularities of their cycles also may be produced the same way as the irregularities of solar cycles. Work on constructing flux transport dynamo models for solar-like stars has just begun \citep{KKC14}. Our ability to model stellar dynamos may throw more light on how the solar dynamo works. \begin{figure} \centering \includegraphics[width=0.9\textwidth]{fig4.eps} \caption{A study of magnetic field evolution on the solar surface from the 3D kinematic dynamo model of \citet{Haz17}, showing how the polar field builds up from a single tilted bipolar sunspot pair due to the Babcock--Leighton mechanism. The different panels show the distribution of magnetic field at the following epochs after the emergence of the bipolar sunspots: (a) 0.025 yr, (b) 0.25 yr, (c) 1.02 yr, (d) 2.03 yr, (e) 3.05 yr, (f) 4.06 yr.} \end{figure} All the theoretical results we discussed are based on axisymmetric 2D kinematic dynamo models. One inherent limitation of such models is that the rise of a magnetic loop due to magnetic buoyancy and the Babcock--Leighton process of generating poloidal flux from it are intrinsically 3D processes and can be included in 2D models only through very crude averaging procedures \citep{Nandy01,Munoz10,CH16}. As we have discussed, the fluctuations in the Babcock--Leighton process play a crucial role in producing the irregularities of the solar cycle. In order to model these fluctuations realistically, it is essential to treat the Babcock--Leighton process itself more realistically than what is possible in 2D models. The next step should be the construction of 3D kinematic dynamo models for which efforts have begun \citep{YM13,MD14,MT16,Haz17}. Such 3D kinematic dynamo models can treat the Babcock--Leighton mechanism more realistically (see Figure~4) and should provide a better understanding of how fluctuations in the Babcock--Leighton mechanism cause irregularities in the dynamo. The magnetic field presumably exists in the form of flux tubes throughout the convection zone and one limitation of a mean field model is that flux tubes are not handled properly \citep{chou03}. A 3D kinematic model allows one to model flux tubes in a more realistic fashion. A proper inclusion of flux tubes in a dynamo model is essential for explaining such interesting observations as the predominance of negative helicity in the norther hemisphere \citep{Pevtsov95}, which is presumably caused by the wrapping of the poloidal field around the rising flux tubes \citep{chou03,Chat06,Hotta11}. This process can be modelled in 2D mean field dynamo models only through drastic simplifications \citep{Chou04}. It should be possible to model this better through 3D kinematic dynamo models. In other words, after the tremendous advances made by the 2D kinematic flux transport model during the last quarter century, it appears that that 3D kinematic dynamo models are likely to occupy the centre stage in the coming years. | 18 | 8 | 1808.08550 |
1808 | 1808.00451.txt | We analyse an N-body simulation of the interaction of the Milky Way (MW) with a Sagittarius-like dSph (Sgr), looking for signatures which may be attributed to its orbital history in the phase space volume around the Sun in light of {\it Gaia} DR2 discoveries. The repeated impacts of Sgr excite coupled vertical and radial oscillations in the disc which qualitatively, and to a large degree quantitatively are able to reproduce many features in the 6D {\it Gaia} DR2 samples, from the median $V_{R},V_{\phi}, V_{z} $ velocity maps to the local $\delta\rho(v_{z},z)$ phase-space spiral which is a manifestation of the global disc response to coupled oscillations within a given volume. The patterns in the {\it large-scale} velocity field are well described by tightly wound spirals and vertical corrugations excited from Sgr's impacts. We show that the last pericentric passage of Sgr resets the formation of the {\it local} present-day $\delta\rho(v_z,z)$ spiral and situate its formation around 500-800 Myr. As expected $\delta\rho(vz,z)$ grows in size and decreases in woundedness as a function of radius in both the {\it Gaia} DR2 data and simulations. This is the first N-body model able to explain so many of the features in the data on different scales. We demonstrate how to use the full extent of the Galactic disc to date perturbations dating from Myr to Gyr, probe the underlying potential and constrain the mass-loss history of Sgr. $\delta\rho(vz,z)$ looks the same in all stellar populations age bins down to the youngest ages which rules out a bar buckling origin. | Gaia's second data release has brilliantly revealed the richness of substructure in our Galactic disc in unprecedented and unforeseen ways \citep{katz18,antoja18}. The existence of substructure in the form of moving groups of stars around the Sun was first pointed out many decades ago and interpreted as a signature of disrupted clusters \citep{eggen69}. Analysis of the {\it Hipparcos} data \citep{dehnen98} offered additional insight that such groups could be excited by asymmetries in the disc, such as the bar or spiral arms \citep{dehnen00, quillen05}. However, other out-of-equilibrium features might be excited through external agents as well, such as mergers, which produce ringing in the $U-V$ plane \citep{minchev09}. These arching features become clearly visible in the space of integrals $E-Lz$, \footnote{These can be loosely understood as integrals on intuitive grounds despite the fact Milky Way potential is non-axisymmetric and time-dependent.} around solar neighbourhood-like (SNhd) regions \citep{gomez12a}. We also note that recent measurements of asymmetries in the $U-V$ plane as a function of height, such as for the Coma Berenices seen in GALAH \citep{quillen18a} and in RAVE \citep{monari18} have also indicated another hint for a merger origin of this moving group. The Sagittarius dwarf galaxy \citep{ibata94}, has long been suggested to be a potential perturber to the Milky Way Galactic disc \citep{ibata98, dehnen98,bailin04}. Indeed, using test particle simulations \cite{quillen09} showed that a Sgr-like satellite could induce perturbations to a disc in the form of a warp, spiral structure as well as streams in the $U-V$ plane. However, only recently has it been possible to assess the impacts of Sagittarius on the Galactic disc with high-resolution live N-body simulations \citep{purcell11,gomez12,gomez13,laporte18a,laporte18b}. These are necessary in order to study the excitation of vertical density waves which are known to be sustained by self-gravitating discs \citep{weinberg1991}, an aspect which cannot be captured by simple phase-mixing models of tracers subject to impulsive perturbations. Direct first evidence of ringing in the Milky Way using the \cite{schuster06} and SEGUE F/G dwarf stellar sample was measured by \cite{gomez12} who attributed the signal to a passage with the Sgr dSph. Co-incidentally, signatures of vertical oscillations in the disc became apparent as local asymmetries in density and velocity seen in spectroscopic surveys \citep{widrow12,williams13,carlin13,schonerich18,carrillo18} and large scale arc-like features above and below the disc plane in the outer disc \citep{newberg02,crane03,grillmair06,martin07,sharma10}. Since then, \cite{antoja18} used the Gaia second data release to uncover a series of substructures in and around the solar neighbourhood, confirming unambiguously that the galaxy is out-of-equilibrium and in a middle of an ongoing phase-mixing process. The authors used a toy model from \cite{minchev09} to estimate the time of impact/disturbance in the solar neighbourhood finding timescales coincident with Sgr's orbital timescales from its last pericentric passage \citep{laporte18b}. This was presented as evidence that the source of perturbation could have been related to Sgr. In this paper, we explore whether any of the various features in the recently reported substructures in velocity space of {\it Gaia} DR2 \citep{antoja18} can be attributed as clear signals to Sgr's orbital history. We do so, using simulations from the suite of numerical N-body experiments in \citep{laporte18b}. These simulations follow the orbit of Sgr-like galaxies around a Milky Way-like host from the first pericentric passage to the present-day. These were the first sets of numerical experiments to successfully reproduce the excitation of the Monoceros Ring and more distant stellar overdensities while producing vertical fluctuations in the solar neighbourhood in good agreement with the constraints set by \cite{widrow12} (see section 2.2 for more details). This paper is organised as follow. In section 2, we describe the simulations used in this contribution. In section 3 we present maps of velocity in regions around the Sun which we compare to the recent results of \cite{katz18} and what these may reflect in the Galaxy. In section 4, we compare models to the recent results of \citep{antoja18}, \cite{monari18} on substructure in the solar neighbourhood. In particular, we look for signs of substructure and phase-mixing which can be related to the Sgr dwarf galaxy. In section 5, we discuss why the Sgr scenario is currently the only one supported by data spanning 10 scale lengths of the disc ruling out alternative models. In section 6, we review various successes of our models and highlight a number of predictions which can be explored with current capabilities beyond Galactic structure/kinematics studies (e.g. through the CMD diagram) as well with future observational campaigns to probe larger volumes of the disc's phase-space. We end the discussion by highlighting a number of ways by which the models could be refined further in the future. | Many of the signs of perturbations in the Milky Way such as the vertical disturbances in the MW within the solar neighbourhood, the spiral patterns in the solar vicinity, ridges in $v_{\phi}-R$, the Monoceros Ring and other overdensities such as EBS, ACS and TriAnd can {\it all be understood as part of the interaction of the disc with the Sgr dwarf galaxy}, with the condition that the progenitor was more massive than previously assumed with $M\sim6\times10^{10} \rm{M_{\odot}}$, a picture that has recently gained support on different but complementary grounds \citep[see][]{deboer14,gibbons16,laporte18b}. In summary, we show that: \begin{enumerate} \item The velocity fields seen in the DR2 data as in \cite{katz18} can be reproduced qualitatively by an interaction of Sgr with the Milky Way disc which excites radial and vertical motions in the disc. These signals get translated in velocity space as a succession of inward and radial ridges in $V_{R}$ as well as kinematic warping of the disc in $V_{Z}$. The radial motion pattern is related to a bi-symmetric spiral which Sgr excites during its last pericentric passage which gets tightly wound at present. This also gives rise to more complex structures such as the ``bridge'' seen in the median radial velocity field. However, the amplitudes in the models considered are larger notably in $V_{R}$ by a factor of $\sim2$. This may be related to the orbital mass loss history of Sgr (which would affect the resulting tides acting on the Galactic disc) and/or the underlying disc model assumed for the MW. \item The interaction with Sgr leaves signs of ongoing phase-mixing around the Sun, with patterns remarkably similar to those recently observed in the $Gaia$ DR2 data \citep{antoja18}. The disagreements are only within a factor of $\sim1.5$ in the $z$ amplitude of the signal, leaving room for more detailed exploration, using different numerical setups specifically tailored to stream fitting Sgr and studying the disc response during the last pericentric passage while varying the Galactic disc structure. \item The interaction model also produces asymmetric moving groups reminiscent of structures seen in the solar neighbourhood such as Coma Berenices (see Appendix). \item Ridges similar to \cite{antoja18} are produced continuously by spiral arms excited by the interaction with Sgr throughout the evolution of the Milky Way. The current patterns seen today in $v_{\phi}-R$ have similar slopes to those excited in the simulations by Sgr, suggesting some possible link between the excitations of the ridges with the interaction with the Sgr. We do not get the same number of observed ridges, but other models of transient spirals do not either \citep{hunt18}. Some prominent ridges can be explained by resonant processes with the bar \citep{fragkoudi19}, thus future explorations will be needed to disentangle which ridges could be associated with a Sgr excitation origin. Indeed the spacing of these ridges every $\sim30 km\,s^{-1}$ in the {\it Gaia} DR2 data bares resemblance with the earlier studies of phase-mixing in the disc in the $U-V$ plane by \cite{minchev09} due to recent mergers. \item We find two previously unknown ridges at $v_\phi\sim120, 90\,\rm{km\,s^{-1}}$ below Arcuturus, confirming predictions of semi-analytic models of Galactic ringing \citep{minchev09}. \item Another prediction from the Sgr impact model is that the phase-space spiral in the Anticenter should harbour signs of past perturbations due to the longer timescales sustained in the disc at large Galactocentric distances, with large amplitudes in the $z$-direction of order $\Delta Z\gg 1 \rm{kpc}$. \item We presented the first map of the phase-space spiral in overdensity $\delta\rho(v_z,z)$, showing clearly more than two wraps (Fig. 7). This would naturally increase the timescales of the onset of the perturbation derived through the use of toy-models by \cite{binney18} which missed this aspect of the data, putting them in better agreement with the accepted orbital timescales for Sgr from stream fitting \citep[e.g.][]{johnston05,law10} of $t\sim0.8\rm{Gyr}$. \item We show that the phase-space spiral has the same shape in all stellar age populations in overdensity $\delta(v_z,z,\tau)$. This means that: 1) all the stars responded the same and 2) that the onset of the perturbation must have been recent, the latter point corroborating \cite{tian18}'s findings. This at face value rules out a bar buckling origin for it in our Galaxy. \item The response of the disc is global and we trace the \cite{antoja18} phase-space spiral for the first time using the DR2 data alone out to appreciably small and larger Galactocentric distances where their shape changes as expected from the underlying potential, a behaviour that is not surprisingly captured in the numerical model. These maps hold crucial information which will be invaluable for any attempts trying to model the zeroth order distribution function of the Milky Way or infer the mass of Sgr during its last pericentric passage. \end{enumerate} We note that \cite{quillen09,purcell11} argued that the Sgr dwarf could be an architect of spiral structure in the Milky Way. At the time, not many constraints around the Sun were available. From our new models, we conclude that the Sgr dwarf may be more closely linked to the disc than previously thought, generating the outer disc structures \citep{laporte18b,laporte18c} all the way to shaping the central part of the Galaxy and the phase-space structure around the the Sun as revealed by the {\it Gaia} satellite. We showed that the impact of Sgr can explain the origins of many different non-equilibrium features, presenting some qualitative/quantitative agreements. Requiring matches between the models and the data (e.g. Sgr stream and disc asymmetries) could tell us a lot about the disc structure and orbital mass-loss history of Sgr. Finally, a number of predictions of the Sgr impact model are awaiting to be confirmed with upcoming spectroscopic surveys which will complement {\it Gaia} which should further open up the potential of dating even more ancient encounters. | 18 | 8 | 1808.00451 |
1808 | 1808.10501_arXiv.txt | The neutral hydrogen~(HI) gas is an important barometer of recent star formation and metal enrichment activities in galaxies. We develop a novel statistical method for predicting the HI-to-stellar mass ratio $\fgas$ of galaxies from their stellar mass and optical colour, and apply it to a volume-limited galaxy sample jointly observed by the Sloan Digital Sky Survey and the Arecibo Legacy Fast ALFA survey. We eliminate the impact of the Malmquist bias against HI-deficient systems on the $\fgas$ predictor by properly accounting for the HI detection probability of each galaxy in the analysis. The best-fitting $\fgas$ predictor, with an estimated scatter of $0.272$ dex, provides excellent description to the observed HI mass function. After defining an HI excess parameter as the deviation of the observed $\fgas$ from the expected value, we confirm that there exists a strong secondary dependence of the mass-metallicity relation on HI excess. By further examining the 2D metallicity distribution on the specific star formation rate vs. HI excess plane, we demonstrate that the metallicity dependence on HI is more fundamental than that on specific star formation rate. In addition, we show that the environmental dependence of HI in the local Universe can be effectively described by the cross-correlation coefficient between HI excess and the red galaxy overdensity $\rho_{cc}{=}-0.18$. This weak anti-correlation also successfully explains the observed dependence of HI clustering on $\fgas$. Our method provides a useful framework for learning HI gas evolution from the synergy between future HI and optical galaxy surveys. | \label{sec:intro} The neutral hydrogen~(HI) gas represents a key intermediate stage in baryon cycling, between the initial accretion from the diffuse circumgalactic or intergalactic medium~\citep{sancisi2008, tumlinson2017} and the formation of dense molecular clouds that directly fuel star formation~\citep{kennicutt2012, lada2012, leroy2013}. The variation of the HI gas reservoir usually precedes the colour transformation of galaxies induced by star formation and quenching~\citep{baldry2004, faber2007}, while regulating the metallicity of the interstellar medium~(ISM) together with galactic outflows~\citep{dalcanton2007, matteucci2012}. In this paper, we develop a statistical framework for connecting the HI gas mass detected by ALFALFA~\citep{haynes2011} to the stellar mass and optical colours of galaxies observed in SDSS~\citep{york2000}, and explore the physical drivers of gas-phase metallicity and the environmental dependence of HI within this framework. As the most important measure of the HI content of a galaxy, the HI-to-stellar mass ratio $\fgas$~(hereafter referred to as HI fraction) has been found to correlate with the optical colour with a scatter of ${\sim}0.4$ dex~\citep{kannappan2004}. Subsequently, \citet{zhang2009} built a photometric estimator of $\fgas$ by introducing an additional scaling of $\fgas$ with the $i$-band surface brightness, reducing the scatter to $0.31$ dex. \citet{li2012} later extended the $\fgas$ estimator by using a linear combination of four parameters~(including stellar mass, stellar surface mass density $\mu_*$, NUV-$r$ colour, and the $g{-}i$ colour gradient), resulting a slightly improved scatter of $0.3$ dex and a more accurate match to the high-$\fgas$ systems observed by ALFALFA. Alternatively, non-linear predictors have been recently developed using machine learning algorithms, which usually require training over a large number of HI-detected systems~\citep{teimoorinia2017, rafieferantsoa2018}. However, current HI surveys like ALFALFA are relatively shallow in depth, and are thus systematically biased against low-$\fgas$ systems at any given redshift. Consequently, any $\fgas$ predictor inferred or trained exclusively from systems above the HI detection threshold would be plagued by the Malmquist bias, overestimating the $\fgas$ for systems that are missed by the HI survey. Such Malmquist bias can be partially alleviated by observing a smaller volume to a higher depth in HI. For example, using a roughly $\fgas$-limited but significantly smaller sample~(GALEX Arecibo SDSS Survey), \citet{catinella2010} constructed a $\fgas$ predictor using the linear combination of NUV-$r$ colour and $\mu$, resulting in a scatter of ${\sim}0.3$ dex~\citep[see also][]{catinella2013}. Without having to trade volume for depth, we develop a new method to eliminate the Malmquist bias when predicting $\fgas$ from the stellar mass and colour of SDSS galaxies, by properly accounting for the ALFALFA detection probability of each SDSS galaxy in the analysis. Beyond $\fgas$, the metal abundance within the gas serves as the fossil record of the chemical enrichment history, reflecting the complex interplay between star formation and gas accretion during the baryon cycling~\citep{peeples2014}. For star-forming galaxies, gas-phase metallicity is tightly corrected with stellar mass with a scatter of $0.1$ dex in the oxygen-to-hydrogen abundance ratio, forming the well-known mass-metallicity relation~\citep[MZR;][]{tremonti2004}. It has been suggested that the star formation rate~(SFR) could drive the scatter in MZR --- there exists a so-called fundamental metallicity relation~(FMR), which manifests as a strong secondary dependence of the MZR on SFR~\citep[][]{mannucci2010, laralopez2010, andrews2013}. Various theoretical models have subsequently been proposed to explain the MZR, assuming SFR is the main process that shaped the MZR~\citep{peeples2011, dave2012, dayal2013, lilly2013, zahid2014}. However, the physical driver of the scatter in MZR is still under debate, and the existence of FMR depends on the systematic uncertainties in the metallicity estimator and potential biases in the sample selection~\citep{yates2012, salim2014, telford2016}. Besides star formation, it is reasonable to expect that gas accretion plays a role in regulating the metallicity of the ISM. Indeed, \citet{bothwell2013} showed that the MZR of ${\sim}4000$ ALFALFA galaxies exhibits a strong secondary dependence on HI mass, with HI-rich galaxies being more metal poor at fixed stellar mass. Applying a principal component analysis over ${\sim}200$ galaxies compiled from several molecular gas surveys, \citet{bothwell2016a} further argued that the underlying driver of MZR is the molecular gas mass, and the FMR is merely a by-product of molecular FMR via the Kennicutt-Schmidt law~\citep{bothwell2016b}. More recently, by stacking the HI spectra of star-forming galaxies along the MZR, \citet{brown2018} confirmed the strong anti-correlation between HI mass and gas-phase metallicity at fixed stellar mass, providing further evidence that the scatter in the MZR is primarily driven by fluctuations in gas accretion. To ascertain whether SFR or HI mass is the more fundamental driver, we will present a comprehensive analysis of metallicity, SFR, and HI mass for a large sample of galaxies jointly observed by SDSS and ALFALFA. In addition to the optical properties of each galaxy, the HI gas reservoir also depends on the large-scale density environment. For example, it is long known that satellite galaxies in massive halos are deficient in HI~\citep{haynes1984, boselli2006, yoon2015, jaffe2015}, due to processes like the ram-pressure and tidal stripping~\citep{gunn1972, merritt1983, moore1996, abadi1999, mccarthy2008, kronberger2008, bekki2009}. Gas accretion history may be tied to the halo growth history, which is known to be correlated with the large-scale environment~\citep{fakhouri2010}. The environmental dependence of cosmic HI distribution can be predicted using semi-analytic models~\citep{fu2010, xie2018} and hydro-dynamic simulations~\citep{dave2017}, or statistically accounted for within the halo model~\citep{guo2017, obuljen2018}. However, a quantitative description of the environmental dependence of HI is still lacking. In our analysis, we quantify this dependence using the cross-correlation coefficient $\rho_{cc}$ between HI excess and galaxy overdensity, and develop three independent approaches to measuring $\rho_{cc}$ directly from data. This paper is organized as follows. We briefly describe the data and the joint SDSS-ALFALFA sample in \ S~\ref{sec:data}, and introduce our likelihood model in \S~\ref{sec:method}. We present our main findings on the mass-metallicity relation in \S~\ref{sec:mzr} and the environmental dependence of HI in \S~\ref{sec:env}. We conclude by summarizing our results and looking to the future in \S~\ref{sec:conc}. Throughout this paper, we assume the {\it WMAP9} cosmology~\citep{wmap2013} for distance calculations. All the length and mass units in this paper are scaled as if the Hubble constant were $100\,\kms\mpc^{-1}$. In particular, all the separations are co-moving distances in units of $\hmpc$, and the stellar and HI mass are both in units of $\hhmsol$. We use $\lg x{=}\log_{10} x$ for the base-$10$ logarithm. | \label{sec:conc} In this paper, we develop a statistical method to infer the HI-to-stellar mass ratio $\fgas$ of galaxies from their stellar mass and optical colour, using a volume-limited galaxy sample jointly observed by SDSS and ALFALFA. Compared to the traditional methods, the key feature of our method is its capability of removing the Malmquist bias against low-$\fgas$ systems in ALFALFA, via a self-consistent modelling of the HI detection rate of each galaxy observed in SDSS. The best-fitting HI fraction predictor has an estimated scatter of $0.272$ dex, slightly smaller than the ${\sim}0.30$ dex reported by traditional methods. To explore the impact of gas accretion on gas-phase metallicity, we define an HI excess parameter $\gamma$ as the deviation of the observed $\lg\,\fgas$ from the expected value (normalized by scatter). We discover a strong secondary dependence of the mass-metallicity relation on $\gamma$, echoing the findings of \citet{bothwell2013} and \citet{brown2018}. This secondary dependence defines a fundamental metallicity relation of HI, similar to the fundamental metallicity relation of the star formation rate~\citep{mannucci2010, laralopez2010, andrews2013}. By taking advantage of two tight scaling relations, i.e., the mass-metallicity relation and the star formation main sequence, we define the relative metallicity and relative $\ssfr$ of each galaxy as the (normalized) deviations from the two respective mean relations. To elucidate the underlying driver of the scatter in the MZR, we examine the 2D relative metallicity distribution on the relative $\ssfr$ vs HI excess plane. We find that the variation of relative metallicity is primarily driven by the change in HI excess, so that galaxies with higher HI excesses always have lower relative metallicities, regardless of the difference in relative $\ssfr$. This 2D metallicity map provides strong evidence that the metallicity dependence on HI is more fundamental that on SFR. Furthermore, the HI excess also depends on the large-scale overdensity environmental. Using the red galaxy overdensity $\dred$ as a measure of the large-scale environment, we demonstrate that there exists a weak anti-correlation between HI excess and $\dred$ in the SDSS-ALFALFA joint sample. From the dependence of detection rate and HI excess on $\dred$, we infer the cross-correlation coefficient $\rho_{cc}$ between the two quantities to be $-0.18$. The $\rho_{cc}{=}-0.18$ model also successfully reproduces the dependence of HI clustering on $\fgas$. We believe this anti-correlation can be largely explained by the ram pressure and tidal stripping of HI gas discs in cluster environments~\citep[but see][]{wang2018}. With the advent of exciting HI surveys like the Square Kilometer Array~\citep[SKA;][]{maartens2015} and the Five-hundred-meter Aperture Spherical Telescope~\citep[FAST;][]{nan2011}, the HI sky will be observed to a much higher depth within a significantly larger volume than ALFALFA. Our method will provide a viable path to the synergy between the next-generation HI surveys and upcoming optical surveys, e.g., the Bright Galaxy Survey program within the Dark Energy Spectroscopic Instrument~\cite[DESI;][]{desi2016}. In particular, we expect the method to provide valuable insight into the evolution of HI gas and metallicity in cluster environments~\citep{peng2014, li2018} and the dependence of HI on large-scale tidal environments~\citep{liao2018, alam2018}. | 18 | 8 | 1808.10501 |
1808 | 1808.02314_arXiv.txt | {It is well-known that the higher policyclic aromatic hydrocarbon (PAH) abundance, lower black hole mass, higher accretion rate and lower luminosities are among the major characteristics of Narrow-Line Seyfert galaxies (NLS1), when they are compared to Broad line Seyfert galaxies (BLS1). NLS1s may be normal Seyfert galaxies at an early stage of evolution, their black holes may still be growing and/or they could be special for some other reason. In this work we discuss the findings that NLS1s have most of line and continuum luminosities correlated with FWHM(H$\beta$), which may be the trace of their rapid black hole mass grow. BLS1 do not show such trends. Also, PAHs may be destroyed as the black hole grows and the starbursts are removed, for NLS1 objects.} \FullConference{Revisiting narrow-line Seyfert 1 galaxies and their place in the Universe - NLS1 Padova\\ 9-13 April 2018 \\ Padova Botanical Garden, Italy} \begin{document} | Narrow-line Seyfert 1 (NLS1) galaxies show relatively narrow full width at half maximum (FWHM$\le$2000 km s$^{-1}$, \cite{Veron01}) of permitted lines, narrower than in typical Seyfert 1 galaxies. NLS1s have certain specific characteristics with respect to the broad line active galactic nuclei (BLAGNs). NLS1 objects have more polycyclic aromatic hydrocarbons (PAHs) \cite{Sani10}, more dust spirals and bars \cite{Deo06} and higher accretion rates \cite{Bian03} than BLAGNs, while BLAGNs have higher black hole (BH) masses (M$_{\rm BH}$) \cite{Mathur00}, and higher optical, X-ray and UV luminosities \cite{Lakicevic18,Grupe10}. NLS1s have lower optical variability \cite{Rakshit17b}, higher X-ray variability \cite{Zhou06} and possibly a lower inclination \cite{Rakshit17a}, than the BLAGNs. NLS1s have strongest Fe\,II emission, strongest X-ray excess, and largest C\,IV blueshifts \cite{Shapovalova12}. It is believed that NLS1s are AGNs in the early stage of evolution \cite{Mathur00,Mathur01} and that their black holes are growing \cite{Jin12a}. The aim of this research was to explore if NLS1s show some characteristics different from BLAGNs, using some optical and mid-infrared (MIR) spectroscopic parameters, as this may suggest the nature of differences between these two groups of objects. Here, we present some correlations that are different for NLS1s and BLAGNs. We found that NLS1s show the dependence of most of line and continuum luminosities and PAH with FWHM(H$\beta$), while BLAGNs do not show these trends. | \begin{itemize} \item The correlations between FWHM(H$\beta$) and luminosities of MIR and optical lines and continuum, for NLS1s, are probably the consequence of the M$_{\rm BH}$ growth which increases both FWHMs and luminosities. Higher mass increase the luminosities and velocities. \item We found that RPAH decreases with M$_{\rm BH}$ and luminosities of coronal lines and continuum, only for NLS1s. \item The trends M$_{\rm BH}$--Luminosities of coronal lines are higher for NLS1s than for BLAGNs. \end{itemize} | 18 | 8 | 1808.02314 |
1808 | 1808.05158_arXiv.txt | We present an extensive study of 72 archival {\it Chandra} light curves of the high-frequency-peaked type blazar Mrk 421, the first strong extragalactic object to be detected at TeV energies. Between 2000 and 2015 Mrk 421 often displayed intraday variability in the 0.3-10.0 keV energy range, as quantified through fractional variability amplitudes that range up to 21.3 per cent. A variability duty cycle of $\sim$ 84 per cent is present in these data. Variability timescales, with values ranging from 5.5 to 30.5 ks, appear to be present in seven of these observations. Discrete correlation function analyses show positive correlations between the soft (0.3--2.0 keV) and hard (2.0--10.0 keV) X-ray energy bands with zero time lags, indicating that very similar electron populations are responsible for the emission of all the X-rays observed by {\it Chandra}. The hardness ratios of this X-ray emission indicate a general ``harder-when-brighter" trend in the spectral behaviour of Mrk 421. Spectral index--flux plots provide model independent indications of the spectral evolution of the source and information on the X-ray emission mechanisms. Brief discussions of theoretical models that are consistent with these observations are given. | \label{sec:introduction} An active galactic nucleus (AGN) involves a supermassive black hole (SMBH), fueled by an accretion disc, producing a variety of highly energetic phenomena \citep{1984ARA&A..22..471R}. When a radio-loud AGN is viewed with one of its relativistic jets in close proximity ($\leq$10$^{\circ}$) to our line of sight, it is categorized as a blazar \citep{1995PASP..107..803U}. Blazars club together BL Lacertae objects, which have nearly featureless optical continua and many flat spectrum radio quasars, (FSRQs) that show extensive broad emission lines \citep[e.g.,][]{ 2015MNRAS.451.3882A}. Blazars are observed to be particularly violent AGNs, involving multiple outstanding attributes, including: dominance of non-thermal emission; high polarization; extreme flux variability across the entire electromagnetic (EM) spectrum; core-dominated radio morphology; and flat radio spectrum. All of these can be understood in terms of relativistic motion of plasma in the jets and Doppler boosting \citep[e.g.,][]{2016MNRAS.455..680A, 2012AJ....143...23G}. The high polarization $( > 3 \%)$ of their radio to optical emission means that the synchrotron emission mechanism is responsible for broadband non-thermal EM radiations from blazars at lower frequencies (radio through the UV or X-ray bands), while at higher frequencies it is likely to be dictated by inverse Compton (IC) scattering of seed photons by the same electrons producing the synchrotron emission. Blazar spectral energy distributions (SEDs) demonstrate a double-peak structure \citep[e.g.,][]{2012AJ....143...23G}. The low energy peaked blazars (LBLs) have the first SED bump peak in mm to optical bands and the second bump at GeV energies, while the high energy peaked blazars (HBLs) have the first component peak at UV/X-ray while the second ranges up to TeV energies \citep{Finke14}. The BL Lac/FSRQ sub-classes also can be distinguished based on optical polarization properties: BL Lac objects show an amplified polarization towards the blue, probably arising due to some intrinsic phenomenon related to the jet-emitting region \citep{1991ApJS...76..813S, 1996MNRAS.281..425M}, while the FSRQs trend in the opposite direction, possibly because of significant contributions from the unpolarized quasi-thermal emission from the accretion disc and surrounding region. \par Blazar observations often show detectable flux variations down to time periods of a few minutes to hours; these must arise from acute physical conditions within small, subparsec scale, regions \citep[e.g.,][]{2016MNRAS.458.1127G}. Blazar variability is conveniently sectioned into three classes, based on their observed time-scales: flux changes occurring over a time-scale of a day or less and up to a few hundredths of a magnitude is termed as intra-day variability (IDV) \citep{1995ARA&A..33..163W}, or microvariability \citep{1989Natur.337..627M}, or intra-night variability \citep{1993MNRAS.262..963G}; variations in flux, typically of a few tenths of a magnitude, that extend from days to weeks are known as short-term variability (STV); while variations ranging from several months to a few years are called long-term variability (LTV) \citep{2004A&A...422..505G}. Extensive studies of STV and LTV for blazars have often shown variations exceeding $\sim$ 1 mag and some have spanned over $\sim$ 5 mag. These flux variabilities in blazars could either be initiated through unstable accretion disc phenomena or solely through changes in the doppler-boosted emission of the relativistic jets \citep{1997ARA&A..35..445U}. Studies of variability timescales and amplitudes serve as key tools in understanding physical processes in the jets and the sizes and locations of the emission regions in AGN. \subsection{\it Mrk 421} Markarian 421 (B2 1101+38; Mrk 421 hereafter) is a nearby elliptical active galaxy ($\alpha_{2000}$ = 11h 04m 27.3139s and $\delta_{2000}$ = +38$^{\circ}$ 12$\arcmin$ 31.7991$\arcsec$) with an intense point-like nucleus, encompassing a $\sim$ 3.6 $\times$ $10^{8}$ M$_{\odot}$ black hole \citep{2008AIPC.1085..399W}. The nuclear source is classified as of the BL Lacertae type as it has a featureless optical spectrum, strongly polarized and variable optical and radio fluxes, and compact radio emission. Mrk 421 has a SED well characterised by a classic two peak shape \citep{1995PASP..107..803U, 1997ARA&A..35..445U}. Most of these observed properties of Mrk 421 are understood to arise from a relativistic jet spotted at a small angle to our line of sight \citep{1995PASP..107..803U}. Relativistic electrons radiating via the synchrotron process produce a non-thermal SED with a polarised continuum extending from the radio to the soft X-ray bands. \par Mrk 421 ($z=0.031$) is one of the closest blazars, at a distance of 134 Mpc (H$_{0} =71$ km s$^{-1}$ Mpc$^{-1}$, $\Omega_{m}$=0.27, $\Omega_{\Lambda}$=0.73) and its synchrotron emission peak was long ago found to lie in the range of 0.1 keV to several keV \citep{1992Natur.358..477P}. The Whipple Cherenkov Telescope claimed to have detected this extragalactic source at TeV energy range (0.5--1.5 Tev) \citep{1993ICRC....1..409S} and it has been confirmed as a TeV source by multiple ground-based $\gamma$-ray telescopes \citep[e.g.][]{2017ApJ...841..100A}. The {\it Compton Gamma-Ray Observatory} (CGRO) easily observed Mrk 421 in the GeV band from space. Mrk 421 is the brightest extragalactic object in $\gamma-$rays in the northern hemisphere \citep[e.g.,][]{2012AJ....143...23G}. \par Thanks to its proximity, observational studies of Mrk 421 are pervasive throughout the entire EM spectrum. The source has had its radio emission followed over the span of 25 years at multiple frequencies \citep[and references therein]{2015MNRAS.448.3121H}. It has shown rapid and extreme optical variability, including LTV of $\sim$ 4.6 mag \citep{1976ARA&A..14..173S}, and IDV up to $\sim$ 1.4 mag of brightness change over a very short period ($\sim$ 2.5 hours) \citep{1988A&AS...72..163X}. Three decades of NIR data reported by \cite{1999ApJS..121..131F} provide IDV and STV confirmation of its blazar nature. In 2006, the source was observed with a peak flux $\sim$ 85 mCrab in the 2.0--10.0 keV band, indicating that the first peak of SED occurred at an energy beyond 10 keV \citep{2009A&A...501..879T, 2009ApJ...699.1964U}. There were reports of ``orphan flares" \citep{2015arXiv150801438F} in TeV $\gamma$-rays, (those not having corresponding increased X-ray emission), in Mrk 421 during 2003 and 2004 multi-wavelength campaigns. On June 10, 2008 Super-AGILE detected a hard X-ray flare. MAXI (Monitor of ALL-sky X-ray Image) marked the strongest X-ray flare in February 2010 ($\sim$ 164$\pm$17 mCrab) \citep{2015ApJ...798...27I}. HESS (High Energy Stereoscopic System) \citep{2005A&A...430..865A} and MAGIC (Major Atmospheric Gamma Imaging Cherenkov Telescopes) \citep{2007ApJ...669..862A} observed the time-average high energy spectrum of Mrk 421 during its flaring stages. \par The dominant synchrotron-self-Compton (SSC) model considers that the same electron population is responsible for the production of soft X-rays and high energy $\gamma$-rays. The SSC model agrees with the results of \cite{2005ApJ...630..130B} and \cite{2009ApJ...695..596H}, where the fluxes were well correlated with a time lag of less than 1.5 days \citep{2016arXiv160509017M}. In April 2013, Mrk 421 was observed to undergo a major X-ray outburst and was comprehensively investigated by multiple observational facilities, including the {\it NuSTAR} and {\it Swift} satellites. Intensive studies of Mrk 421 probing its multi-wavelength (MW) behaviour are numerous \citep[and references therein]{1994IAUC.5993....2T, 1995ApJ...438L..59K, 2004A&A...422..505G, 2008ATel.1574....1C, 2008ATel.1583....1P, 2008ICRC....3..973S, 2012A&A...545A.117L, 2012MNRAS.420.3147G, 2013A&A...559A..75B, 2014A&A...570A..77P, 2015ApJ...811..143P, 2015A&A...580A.100S, 2015MNRAS.448.3121H, 2016ApJ...819..156B}. MW campaigns incorporating {\it Fermi-LAT} gamma-ray detections comprehensively studied Mrk 421 and produced its first ever complete $\gamma$-ray continuum during a quiescent state \citep{2014ApJ...782..110A}. A multi-decade optical light curve spanning 1900 to 1991 was extracted by \cite{1997A&AS..123..569L} in the B-band, suggesting two possible observed time periods of 23.1$\pm$1.1 yrs and 15.3$\pm$0.7 yrs in those flux variations \citep{2016A&A...591A..83S}. A significant amount of correlation ($\sim$ 68$\%$) was found with X-ray data from {\it RXTE-ASM}, when \cite{2010tsra.confE.197T} studied the long term VHE light curve of the source. Strong episodes of TeV--Xray correlation were discussed by \cite{2005A&A...433..479K}. The typical nature of moderate X-ray--GeV flux correlations has been recently examined by \cite{2015ApJ...806...20B} through multi-wavelength observations made from 2008 to 2013. \par The least well understood aspect of blazar variability is probably seen on IDV timescales. To search for and analyze IDV in blazars, we are working on a project in which we study data taken with various ground and space based telescopes \citep{2008AJ....135.1384G, 2008AJ....136.2359G, 2012MNRAS.425.1357G, 2016MNRAS.458.1127G, 2017MNRAS.465.4423G, 2010ApJ...718..279G, 2012AJ....143...23G, 2012MNRAS.420.3147G, 2012MNRAS.425.3002G, 2015MNRAS.452.4263G, 2012MNRAS.424.2625B, 2015MNRAS.450..541A, 2015MNRAS.451.3882A, 2016MNRAS.455..680A, 2015MNRAS.451.1356K, 2017ApJ...841..123P}. In this paper, we present a study of the IDV of Mrk 421 study using the {\it Chandrasekhar X-ray Observatory} satellite. We employ all the archival data taken by {\it Chandra} since its launch, extending from 2000 May 29 to 2015 July 02 ($\approx$ 16 years) and totaling 72 IDV light curves.This is the most extensive IDV study of Mrk 421 in the X-ray band, covering the longest temporal span. This work provides us with better understanding of the X-ray variability properties of Mrk 421, along with the correlations between hard and soft X-ray bands. \begin{table*} Table 1. Observation log of {\it Chandra} data for Mrk 421. \centering \noindent \scalebox{1.2}{ \begin{tabular}{rcccccccc} \hline \hline ObsID & Date of Observation & Start Time (UT) & Detector & Grating & Exposure Time \\ & (dd-mm-yyyy) & (hh:mm:ss) & & & (ks) \\ \hline 1714 & 29-05-2000 & 11:39:48 & ACIS-S & HETG & 19.83 \\ 1715 & 29-05-2000 & 17:40:11 & HRC-S & LETG & 19.84 \\ 4148 & 26-10-2002 & 00:05:02 & ACIS-S & LETG & 96.84 \\ 4149 & 01-07-2003 & 14:19:35 & HRC-S & LETG & 99.98 \\ 5318 & 06-05-2004 & 14:53:40 & ACIS-S & LETG & 30.16 \\ 5171 & 13-07-2004 & 17:19:41 & ACIS-S & LETG & 67.15 \\ 5332 & 14-07-2004 & 12:21:48 & ACIS-S & LETG & 67.06 \\ 8378 & 07-01-2007 & 23:08:15 & ACIS-S & LETG & 28.16 \\ 6925 & 08-01-2007 & 07:12:48 & ACIS-S & LETG & 27.69 \\ 8396 & 21-01-2007 & 07:52:59 & HRC-S & LETG & 29.48 \\ 10671 & 10-10-2009 & 23:53:51 & ACIS-S & LETG & 28.96 \\ 10664 & 08-11-2009 & 05:00:56 & ACIS-S & LETG & 20.06 \\ 11605 & 16-11-2009 & 03:43:14 & ACIS-S & LETG & ~5.24 \\ 11606 & 18-11-2009 & 19:16:47 & ACIS-S & LETG & ~5.24 \\ 11607 & 22-11-2009 & 09:17:23 & ACIS-S & LETG & ~5.14 \\ 11960 & 02-02-2010 & 03:31:10 & ACIS-I & LETG & 20.13 \\ 11961 & 04-02-2010 & 08:50:00 & ACIS-I & LETG & 20.17 \\ 11962 & 04-02-2010 & 14:44:43 & ACIS-I & LETG & 19.77 \\ 11963 & 06-02-2010 & 08:54:32 & ACIS-I & LETG & 20.18 \\ 11964 & 06-02-2010 & 14:56:02 & ACIS-I & LETG & 19.77 \\ 11967 & 06-02-2010 & 20:37:52 & ACIS-I & LETG & 19.78 \\ 10663 & 13-03-2010 & 02:11:06 & ACIS-S & HETG & 15.13 \\ 11970 & 13-03-2010 & 06:56:51 & ACIS-S & LETG & 10.07 \\ 10665 & 13-03-2010 & 10:08:56 & HRC-S & LETG & 10.15 \\ 12121 & 13-03-2010 & 13:07:43 & ACIS-S & LETG & 10.07 \\ 10667 & 13-03-2010 & 16:14:18 & HRC-S & LETG & 10.18 \\ 10668 & 13-03-2010 & 19:13:05 & ACIS-S & LETG & 10.06 \\ 10669 & 13-03-2010 & 22:19:39 & HRC-S & LETG & 10.18 \\ 11966 & 14-03-2010 & 01:18:27 & ACIS-S & LETG & 30.06 \\ 10670 & 14-03-2010 & 10:00:29 & ACIS-S & HETG & 14.80 \\ 12122 & 10-07-2010 & 18:37:09 & HRC-S & LETG & 25.18 \\ 13097 & 16-02-2011 & 18:33:05 & ACIS-S & LETG & 30.06 \\ 13098 & 04-07-2011 & 02:15:37 & ACIS-S & HETG & 14.80 \\ 13099 & 04-07-2011 & 06:59:47 & ACIS-S & LETG & 10.06 \\ 13100 & 04-07-2011 & 10:12:04 & HRC-S & LETG & 10.15 \\ \hline \end{tabular}} \end{table*} \begin{table*} Table 1. continued. \noindent \scalebox{1.2}{ \begin{tabular}{lcccccccc} \hline \hline ObsID & Date of Observation & Start Time (UT) & Detector & Grating & Exposure Time \\ & (dd-mm-yyyy) & (hh:mm:ss) & & & (ks) \\ \hline 13104 & 04-07-2011 & 22:07:18 & HRC-S & LETG & 10.18 \\ 13105 & 05-07-2011 & 01:06:06 & ACIS-S & HETG & 15.00 \\ 14266 & 07-04-2012 & 20:29:18 & ACIS-S & LETG & 30.05 \\ 14320 & 03-07-2012 & 11:34:05 & ACIS-S & HETG & 15.03 \\ 14322 & 03-07-2012 & 16:07:49 & HRC-S & LETG & 10.01 \\ 14396 & 03-07-2012 & 19:09:24 & HRC-S & LETG & ~9.79 \\ 14321 & 03-07-2012 & 22:06:42 & ACIS-S & LETG & 10.05 \\ 14323 & 04-07-2012 & 01:13:05 & HRC-S & LETG & 10.18 \\ 14324 & 04-07-2012 & 04:09:45 & HRC-S & LETG & ~9.79 \\ 14325 & 04-07-2012 & 07:07:03 & ACIS-S & LETG & 10.05 \\ 14326 & 04-07-2012 & 10:13:26 & HRC-S & LETG & 10.19 \\ 14397 & 04-07-2012 & 13:10:06 & HRC-S & LETG & ~9.79 \\ 14327 & 04-07-2012 & 16:07:26 & ACIS-S & HETG & 15.04 \\ 15607 & 07-02-2013 & 19:15:19 & ACIS-S & LETG & 30.07 \\ 15476 & 03-04-2013 & 00:36:24 & ACIS-S & LETG & 30.05 \\ 15477 & 30-06-2013 & 16:15:37 & ACIS-S & HETG & 14.65 \\ 15478 & 30-06-2013 & 20:46:21 & ACIS-S & LETG & 10.06 \\ 15479 & 30-06-2013 & 23:58:37 & HRC-S & LETG & 10.14 \\ 15480 & 01-07-2013 & 02:57:24 & ACIS-S & LETG & ~9.76 \\ 15481 & 01-07-2013 & 05:56:13 & HRC-S & LETG & 10.19 \\ 15482 & 01-07-2013 & 08:55:01 & ACIS-S & LETG & ~9.77 \\ 15483 & 01-07-2013 & 11:53:49 & HRC-S & LETG & 10.19 \\ 15484 & 01-07-2013 & 14:52:38 & ACIS-S & HETG & 14.50 \\ 16474 & 06-03-2014 & 08:13:12 & ACIS-S & LETG & 60.07 \\ 16424 & 25-06-2014 & 13:54:38 & ACIS-S & HETG & 15.05 \\ 16425 & 25-06-2014 & 20:44:18 & ACIS-S & LETG & 10.08 \\ 16426 & 25-06-2014 & 23:56:09 & HRC-S & LETG & 10.15 \\ 16427 & 26-06-2014 & 02:55:01 & ACIS-S & LETG & 10.08 \\ 16428 & 26-06-2014 & 06:01:20 & HRC-S & LETG & 10.18 \\ 16429 & 26-06-2014 & 09:00:09 & ACIS-S & LETG & 10.08 \\ 16430 & 26-06-2014 & 12:06:31 & HRC-S & LETG & 10.17 \\ 16431 & 26-06-2014 & 15:05:17 & ACIS-S & HETG & 15.05 \\ 17385 & 01-07-2015 & 00:42:11 & ACIS-S & HETG & 15.03 \\ 17387 & 01-07-2015 & 20:54:30 & HRC-S & LETG & 10.18 \\ 17389 & 01-07-2015 & 23:51:09 & HRC-S & LETG & 10.19 \\ 17391 & 02-07-2015 & 02:46:19 & HRC-S & LETG & 10.18 \\ 17392 & 02-07-2015 & 05:43:39 & ACIS-S & HETG & 14.06 \\ \hline \end{tabular}} \end{table*} \par The paper is organized as follows. Section 2 briefly describes the {\it Chandra} satellite instrumentation along with the methodology for data reduction. Data analysis techniques used to search for flux and spectral variability are discussed in Section 3. Section 4 and Section 5 give our results and a discussion, respectively. Our conclusions are reported in Section 6. | We studied 72 {\it Chandra} light curves of the TeV blazar Mrk 421, and searched for variability timescales of IDV. The rapid X-ray variability studied here most likely originates within compact regions of the relativistic jet. Our conclusions are summarised as follows: \begin{enumerate} \item The fractional variability amplitude provides a clear indications of variability on many occasions, with highest variability amplitude being over 21 per cent. The variability in hard energy X-ray bands presumably originates from a compact region within a relativistic jet. \item The duty cycle for these variations is at least $\approx$ 84 per cent which indicates that the source was exceptionally variable in the observed 16 year span. \item We found evidence for timescales ranging from 5.5 to 30.5 ksec in 7 LCs of Mrk 421 using the ACF technique. Other observations have noisier ACF plots in which variability timescales are not clearly present. Using the shortest strong variability timescale of 9.51 ks, we can estimate key parameters in a fashion this is essentially independent of the theoretical model. We find a magnetic field $B \textgreater 0.10 ~\nu^{-1/3}_{18}$ G, electron Lorentz factor $\gamma \geq 3.06 \times10^{5} \nu^{2/3}_{18}$ and radius of the emitting region $R \leq 6.92 \times 10^{15}$ cm. \item The DCF technique was applied to the hard (2--10 keV) and soft (0.3--2 keV) X-ray bands and displayed positive correlations with no time lag. This implies that the emission in both nearby bands arose from the same production region at the same time; i.e., there is no evidence that the softer X-rays arise from synchrotron emission while the harder come from SSC. \item A hardness ratio analysis was also employed to study spectral variations. HRs showed flatter spectra at high fluxes (Fig.\ 2). This indicated that the HR normally increased with increasing flux and got ``harder-when-brighter". Fig.\ 5 displays more information about the spectral evolution of the source. During these extended observations, a few epochs had hard-lags, pointing to the particle acceleration mechanism being responsible for the X-ray emissions, whereas a few epochs had soft-lags, indicating that X-rays are predominantly emitted by the synchrotron cooling mechanism during those periods. \end{enumerate} | 18 | 8 | 1808.05158 |
1808 | 1808.10447_arXiv.txt | Tidal disruption events of stars by supermassive black holes have so far been discovered months to years after the fact. In this paper we explore the short, faint and hard burst of radiation is emitted at maximum compression, as a result of shock breakout. The detection of this burst can be used to capture tidal disruption events in real time. We verify that shock breakout from main sequence stars produces radiation in the X-ray range, but find that it is difficult to detect using all sky X-ray surveying telescopes. In the case of shock breakout from red giants, most of the radiation is emitted in the UV and visible range, which is significantly easier to detect. A similar burst of UV/optical radiation will also be emitted by stars puffed by tidal heating from a previous passage close to the central black hole. This radiation can be detected by surveys like ZTF and LSST. We calculate detection rates for different types of galactic nuclei. For the case of a very full loss cone we predict a detection rate of once per month with LSST, whereas for the case of a very empty loss cone we predict a rate of once per year with LSST. Evidence from a recent tidal disruption event, ASASSN-14li, seems to favour a very full loss cone, in which case LSST is expected to detect one such event every month. | A star that passes too close to a black hole will be ripped apart by the tidal forces exerted by the black hole. Tidal disruption events occur in galactic nuclei \citep{rees_1988}. Recently, a tidal disruption event candidate has even been detected in a globular cluster, suggesting a possible interaction with an intermediate mass black hole \citep{lin_stradler_et_al_2018}. As will be explained in detail in the next section, tidal disruption begins when the distance between the black hole and the star drops below a critical distance called the tidal radius. The time the star spends inside the tidal sphere (i.e. a sphere with tidal radius around the black hole) is of the order of the dynamical time of the star (hours for main sequence stars). Yet, radiation is produced on much longer time scales, ranging from months to years depending on the mass of the black hole \citep{guillochon_ramirez_ruiz_2015}. This means that all tidal disruption events observed so far were captured a months to years after the fact. In this paper we discuss an emission of radiation that occurs while the star is still inside the tidal sphere. This emission is due to the radiative shock breakout due to the compression of the star in a direction normal to the plane of motion. The formation of a shock around the equatorial plane and its motion toward the poles has been shown in hydrodynamic simulation simulations \citep{1983A&A...121...97C, guillochon_et_al_2009}. The actual shock breakout occurs on a very thin layer close to stellar surface, and involves complicated interaction between the gas and the radiation field, so it cannot be analysed using the same numerical simulation. In this work we apply the formalism for shock breakout, which has been studied extensively for supernova explosions \citep{nakar_sari_newtonian_breakout_2010,katz_et_al_2010}, to calculate the luminosity, duration and spectra from radiative shock breakouts from tidal disruption events. The plan of the paper is as follows, in Section \ref{sec:analytic} we derive the formulae for the luminosity, duration and light - curves for radiative shock breakouts from tidal disruption events. In section \ref{sec:observability} we give predictions for the rate at which we expect these events to be detected using present day and next generation instruments. Finally, we conclude in Section \ref{sec:conclusions}. | \label{sec:conclusions} In this work we considered the emission from a radiative shock breakout close to maximum compression, in the course of a tidal disruption of a star by a black hole. We considered two types of stars: main sequence stars and puffed stars. The puffed stars include red giants and main sequence stars puffed by tidal heating from the black hole. For main sequence stars we expect a luminosity of about $10^{41}$ erg/s, a duration of about 20 seconds and an average photon energy of about 1 - 10 keV. In the case of puffed stars, the luminosity is about $10^{40}$ erg/s, a duration of a few hours and an average energy of a few dozen eV. Since the emission from main sequence stars is mostly in the X-ray range, and due to sensitivity of all sky X-ray surveys, we do not expect any such event will be discovered. On the other hand, the emission from puffed stars is emitted in the UV and visible range, where the instruments are more sensitive. We have considered several scenarios for the conditions at other galactic nuclei, and calculated the detection rates in each. In the most optimistic scenario we predict one discovery per century with ZTF. For LSST, our predictions range from once per year in the most optimistic case, to once per century in the pessimistic case. Circumstantial evidence from the tidal disruption event ASASSN-14li seems to favour the most optimistic scenario. One channel that can boost the detection rates, but has not been considered in the previous sections is tidal disruption events from intermediate mass black holes in globular clusters \citep{2018arXiv180608385F}. It has been suggested that they occur at a comparable rate to tidal disruption events in galactic nuclei. Compared to a super - massive black hole, the radiative shock breakout from an intermediate mass black hole would produce a dimmer, but more prolonged emission. This trade - off is due to the fact that the total energy is independent of the mass of the black hole. The detection rate scales linearly with the duration, but super - linearly with the luminosity ($\propto L^{3/2}$ where $L$ is the luminosity), so even if tidal disruption occur at the same in globular clusters and galactic nuclei, it is the latter kind that would dominate the sample of observed shock breakouts. Finally, we want to point out that very little is known about nuclei of other quiescent galaxies. Even a non detection of radiative shock breakouts within a certain period will tell us something about them. For example, a non detection within a decade with LSST will suggest that some process, like collisions with stellar mass black holes, destroys puffed stars. | 18 | 8 | 1808.10447 |
1808 | 1808.09189_arXiv.txt | Using the PMO-13.7m millimeter telescope at Delingha in China, we have conducted a large-scale simultaneous survey of $^{12}$CO, $^{13}$CO, and C$^{18}$O J=1-0 emission toward the Rosette molecular cloud (RMC) region with a sky coverage of $3.5\arcdeg \times 2.5\arcdeg$. The majority of the emission in the region comes from the RMC complex with velocities lying in the range from -2 km s$^{-1}$ to 20.5 km s$^{-1}$. Beyond this velocity range, 73 molecular clumps are identified with kinematic distances from 2.4 kpc to 11 kpc. Based on the spatial and velocity distribution, nine individual clouds, C1-C9, have been identified for the RMC complex. It appears that the C3 cloud is different from other clouds in the RMC complex in view of its characteristic velocity, excitation temperature, and velocity dispersion. Most of the young stellar clusters in the region are located in positions of both high column density and high excitation temperature. Seven new molecular filaments are discovered in the RMC complex. Evidence for cloud-cloud collision is found in the region of young stellar clusters REFL9 and PouF, showing that these young stellar clusters probably result from a cloud-cloud collision. The abundance ratios of $^{13}$CO to C$^{18}$O in the region have a mean value of 13.7 which is 2.5 times larger than the solar system value, showing that UV photons from the nearby OB clusters have strong influence on the chemistry of clouds in the RMC complex. | \label{sec:intro} Stars form in cold and dense molecular clouds. Large scale surveys of molecular clouds in the Milky Way revealed that the major part of molecular clouds in the Milky Way is accumulated into cloud complexes which are called Giant Molecular Clouds (GMCs) \citep{1987ApJ...322..706D,2001ApJ...547..792D}. GMCs are found to be in virial equilibrium state and bound by gravity while the constituent clouds of GMCs and isolated molecular clouds with M $< 10^3$ M$_\sun$ are not in self-gravitational equilibrium \citep{1987ApJ...319..730S, 2001ApJ...551..852H}. The mass function of molecular clouds follows a power law with an index of around -1.7 \citep{1997ApJ...476..166W, 2005PASP..117.1403R} and there exists a scaling relation between the line-width and the size of molecular clouds \citep{1981MNRAS.194..809L, 1987ApJ...319..730S}. It has long been realized that GMCs have complex and hierarchical structure which can be divided into substructures of clouds, clumps, and cores \citep{1999ASIC..540....3B}. With high spatial resolution and sensitivity in the sub-millimeter regime, Herschel observations reveal the ubiquitous presence of filaments in molecular clouds \citep{2014prpl.conf...27A} and have found that star formation in molecular clouds occurs mainly at the junctions of filaments \citep{2012A&A...540L..11S}. Turbulence has been proposed to be responsible for the origin of the hierarchical and filamentary structure of molecular clouds \citep{1994ApJ...423..681V, 2001ApJ...553..227P}. Stellar feedbacks, in particular that from young OB clusters, have strong influence on the evolution of molecular clouds through stellar winds, outflows, and radiation \citep{2002ApJ...566..302M, 2017ApJ...850..112R}. Recently, both numerical simulations and observations have shown that cloud-cloud collision may play an important role in the dynamics of molecular clouds and may be an important mechanism for triggered star formation \citep{2009ApJ...696L.115F, 2010MNRAS.405.1431A, 2017ApJ...835L..14G}. However, it remains unclear what are roles of turbulence, stellar feedbacks, and cloud-cloud collision in the formation, evolution, and destruction of molecular clouds. Observations of GMCs in varying environments with large spatial dynamic range is essential to address this important question. The Rosette molecular cloud (RMC) is an ideal target to analyse the internal structure of GMCs and to characterize the influence of stellar feedbacks. It is associated with an optical emission nebula, the Rosette Nebula. The Rosette Nebula is a well known \ion{H}{2} region which is driven by the NGC 2244 OB cluster. This OB cluster contains six massive stars of O spectral type that have a total luminosity of $\sim$ 10$^6$ L$_{\odot}$ \citep{1985A&A...144..171C}. The expanding \ion{H}{2} region interacts with the surrounding RMC and the photons from the NGC 2244 OB cluster produce a photon dominated region (PDR) at the interface between the \ion{H}{2}-region and RMC \citep{1998A&A...338..262S}. The distance to the young NGC 2244 OB cluster has been estimated using stellar photometry with results ranging from 1.4-1.7 kpc \citep{1981PASJ...33..149O,2002AJ....123..892P}. Using optical spectroscopy \cite{2000A&A...358..553H} determined the orbital and fundamental stellar parameters for each of the components of the V578 Mon binary, which is a member of the NGC 2244 OB cluster, and derived a distance of 1.39 $\pm$ 0.1 kpc for the NGC 2244 OB cluster. As in \cite{2009MNRAS.395.1805D}, we adopted a distance of 1.4 kpc for RMC in this work. The RMC has been surveyed in multi-wavelength. \cite{1997ApJ...477..176P} imaged a region of ~0.7 deg$^2$ toward the RMC in near-infrared (JHK) and detected seven young embedded clusters, PL01-PL07. The FLAMINGOS survey conducted by \cite{2008ApJ...672..861R} confirmed the existence of PL01-P07 and detected four more young clusters, REFL08-REFL10 and the NGC 2237 cluster. The Spitzer telescope survey of the RMC with IRAC and MIPS, covering a region of 1\degr $\times$ 1.5\degr, identified a total of 751 young stellar objects with infrared excess down to a mass limit of 0.4 M$_\sun$. The observations confirmed the seven clusters of \cite{1997ApJ...477..176P} and the existence of clusters REFL08 and REFL09. Furthermore, two new small clusters, PouC and PouD, were detected near clusters PL02 and PL03, respectively \citep{2008MNRAS.384.1249P}. Herschel observations with the PACS and SPIRE instruments in the wavelength range of 70-520 $\mu$m have revealed protostars, gas clumps, filaments, and dust temperature distribution in the RMC \citep{2010A&A...518L..84H, 2010A&A...518L..91D, 2010A&A...518L..83S, 2012A&A...540L..11S}. The first large scale CO J = 1-0 emission survey of the RMC was made with the 1.2 m telescope at Columbia University in New York \citep{1980ApJ...241..676B}. The survey revealed that the RMC is extended along the galactic plane with a maximum extent of $\sim$ 100 pc and it contains emission maxima IRS and A-J. They found that the RMC possesses an overall velocity gradient of 0.20 km s$^{-1}$ pc$^{-1}$. \citet{1986ApJ...300L..89B} mapped the RMC in J=1-0 lines of CO and $^{13}$CO with the 7 m telescope at Bell Laboratories and found that molecular clumps in the RMC are embedded in low volume density interclump molecular gas. By analyzing the 7 m telescope survey data, \citet{1995ApJ...451..252W} found an overall velocity gradient of 0.08 km s$^{-1}$ pc$^{-1}$ for the RMC and interpreted this velocity gradient as the rotation of the RMC. Using the $^{12}$CO and C$^{13}$O J = 1-0 survey data obtained with the 14 m telescope of the Five College Radio Astronomy Observatory, \citet{2006ApJ...643..956H} examined the role of turbulent fragmentation in regulating the efficiency of star formation and it was found that the effect of turbulent fragmentation must be limited and nonexclusive in the RMC. \citet{2009MNRAS.395.1805D} conducted a large-scale survey of the $^{12}$CO J=1-0 emission covering 4.8 deg$^2$ with the James Clerk Maxwell Telescope (JCMT) and it was shown that the dominant bulk molecular gas motion in the region is expansion away from the O stars in NGC 2244. In the present work we present a new large-scale (3.5\arcdeg $\times$ 2.5\arcdeg) survey of $^{12}$CO, $^{13}$CO, and C$^{18}$O J=1-0 emission toward the RMC region. The survey is described in Section 2 and the results are presented in Section 3. We discuss our results in Section 4 and present the summary in Section 5. | \subsection{Distribution of velocity, excitation temperature, and velocity dispersion of RMC}\label{discussion:1} The velocity distribution of molecular clouds in the whole RMC complex is presented in Figure \ref{fig15}. As the C3 cloud spatially overlaps with clouds C2 and C5, its velocity distribution is presented separately in the left panel of Figure \ref{fig15} while the velocity distribution for other clouds in the region, i.e. C1-C2 and C4-C9, is presented in the bottom panel. From Figure \ref{fig15} we can see that clouds of relatively red-shifted velocities are distributed to the southwest of the NGC 2244 cluster while the relatively blue-shifted velocity clouds, in particular clouds C2 and C3, are located to the northeast of the NGC 2244 cluster. This velocity distribution has lead \cite{1995ApJ...451..252W} to conclude that the RMC as a whole possesses a large-scale rotation. However we note that the velocity distribution as shown in Figure \ref{fig15} does not fit well with the rotation scenario, for example, clouds C5 and C6 are located to the left of the dashed line and they have similar velocities to clouds C7-C9 on the right side. Therefore, we propose that no large-scale rotation exists for RMC complex. \begin{figure}[h] \centering \includegraphics[width=0.4\textwidth]{fig15a-group1_m1.pdf} \includegraphics[width=0.4\textwidth]{fig15b-group2_m1.pdf} \caption{Left: velocity distribution of cloud C3 from $^{13}$CO emission. Right: velocity distribution from $^{13}$CO emission for other clouds in the region, i.e., C1-C2 and C4-C9. The dash line shows the rotation axis from \cite{1995ApJ...451..252W}.} \label{fig15} \end{figure} Figure \ref{fig16} shows the excitation temperature distribution of the RMC complex. High temperature ($\sim$ 25 K) occurs at the interface between the Rosette Nebula and the surrounding molecular clouds. The highest temperature (37 K) occurs at the western part of cloud C5. Away from the interaction interface the excitation temperature gradually decreases to a value of around 10-15 K. This temperature distribution is consistent with the Hershel result on dust temperature in the region \citep{2010A&A...518L..83S}. This temperature gradient was first found by \citet{1990A&A...230..181C} from analysis of IRAS images. They showed that the outer part of RMC has stronger IRAS 60 $\mu$m and 100 $\mu$m emission than the interaction interface while the IRAS 12 $\mu$m emission is strong around and inside the \ion{H}{2} region. We note that although cloud C3 is a major cloud in the RMC region, no excitation temperature higher than 20 K is found in this cloud. For a comparison, clouds C5 and C6 are located at a similar distance from the NGC 2244 cluster, but both of them possess temperatures higher than cloud C3. It appears that cloud 3 has not been influenced by the NGC 2244 cluster in view of its low excitation temperature. \begin{figure}[h] \centering \includegraphics[width=0.4\textwidth]{fig16-Tex.pdf} \caption{Distribution of excitation temperature of the RMC complex. The contours outline $^{13}$CO emission region of clouds C1-C2 and C4-C9.} \label{fig16} \end{figure} The velocity dispersion of the RMC complex is shown in Figure \ref{fig17}. As in Figure \ref{fig3}, the circles, squares, and triangles indicate the locations of embedded young stellar clusters in the region. From Figure \ref{fig17} we can see that most of the regions of high velocity dispersion in cloud C3 are associated with embedded young stellar clusters, showing that feedbacks from young stellar clusters are the main cause for the high velocity dispersion. On the other hand, we find that for clouds C1-C2 and C4-C9 the regions of high velocity dispersion are mainly distributed along the \ion{H}{2} region interface. In particular, as we discussed previously in Section 3, the velocity dispersion in cloud C8 is unusually strong and may be caused by the combined influences from the NGC 2244 and NGC 2237 OB clusters. From above discussion it appears that cloud C3 is different from other clouds in the RMC complex in view of its velocity, its temperature, and velocity dispersion distribution. Therefore, we propose that the molecular clouds in the RMC complex can be generally divided into two groups, with cloud C3 as one group (group 1) that shows little impacts from the NGC 2244 and NGC 2237 clusters while other clouds, C1-C2 and C4-C9, as another group (group 2) that shows apparent influence from the NGC 2244 and/or NGC 2237 clusters. We present the column density distribution for group 1 and group 2 clouds separately in Figure \ref{fig18}. We have estimated the distances to the group 1 and group 2 clouds using stellar extinction data. Based on 5-band grizy Pan-STARRS 1 (PS1) and 3-band 2MASS photometry, \citet{2015ApJ...810...25G} trace the extinction on 7$\arcmin$ scales and have obtained a three-dimensional map of interstellar dust reddening. We select 3 regions from group 1 and 4 regions from group 2 with a radius of 7$\arcmin$. There is only one velocity component in these regions. From the map we find a rapid increase in extinction centered at the distance modulus of 11 (1.4 kpc) for both groups of the clouds. Therefore, the both groups of the clouds are located at the same distance as the Rosette Nebula. The age of the NGC 2244 cluster is estimated to be approximately $4 \times 10^6$ yr while that of the Rosette Nebula is about one order of magnitude younger \citep{1981PASJ...33..149O}. The age of the NGC 2237 cluster is 2 $\times$ 10$^6$ yr \citep{2010ApJ...716..474W}. The nearest known supernova remnant (SNR) to the RMC complex is SNR G205.5+0.5 (Monoceros Nebula) which is located about 2.6$\arcdeg$ to the northwest and has an angular diameter of 220$\arcmin$ \citep{2014BASI...42...47G}. With $^{12}$CO and $^{13}$CO data, \citet{2017ApJ...836..211S} identified six positions, denoted as a-f in their Figure \ref{fig1}, where they confirmed that SNR G205.5+0.5 is interacting with the surrounding molecular clouds. However, these positions are well separated from the RMC complex. Therefore, we propose that the stellar feedback to the RMC complex is dominated by the effects of photoionisation and winds from the OB clusters NGC 2244 and NGC 2237, although the nearby supernova remnant G205.5+0.5 may also play a role. \begin{figure}[h] \centering \includegraphics[width=0.4\textwidth]{fig17a-group1_m2.pdf} \includegraphics[width=0.4\textwidth]{fig17b-group2_m2.pdf} \caption{Left: velocity dispersion of cloud C3 in $^{13}$CO emission. Right: velocity dispersion in $^{13}$CO for clouds C1-C2 and C4-C9. The dashed circle indicates the range of influence of the NGC 2244 OB cluster. The circles, squares, and triangles indicate the locations of embedded young stellar clusters as in Figure \ref{fig3}.} \label{fig17} \end{figure} \begin{figure}[h] \centering \includegraphics[width=0.4\textwidth]{fig18a-N_H2_group1.pdf} \includegraphics[width=0.4\textwidth]{fig18b-N_H2_group2.pdf} \caption{Left: column density distribution of group 1 cloud (C3). Right: column density distribution of group 2 clouds (C1-C2 and C4-C9).} \label{fig18} \end{figure} \begin{figure}[h] \centering \includegraphics[width=0.4\textwidth]{fig19-13_18.pdf} \caption{Abundance ratio of $^{13}$CO to C$^{18}$O as a function of the projection distance from the NGC 2244 OB cluster for clouds in group 1 (blue circles) and clouds in group 2 (red triangles). The green line indicates the solar system value (5.5).} \label{fig19} \end{figure} \subsection{Isotopologue abundance ratio } Ultraviolet (UV) photons can dissociate CO and its isotopologues, therefore, have strong influence on the chemistry of molecular clouds. Due to the large difference in optical depth, $^{13}$CO and C$^{18}$O molecules are selectively dissociated by UV photons, with the $^{13}$CO abundance being less affected than the C$^{18}$O abundance. The abundance ratio between $^{13}$CO and C$^{18}$O, R$_{13,18}$, for molecular clouds that are irradiated by massive stars is expected to exceed the values for molecular clouds without strong UV irradiation. To investigate the effects of UV photons from the NGC 2244 and NGC 2237 OB clusters on the RMC complex, we calculated R$_{13,18}$ for clouds in both group 1 and group 2, and the results are presented in Figure \ref{fig19}. To ensure a high reliability of C$^{18}$O detection, only pixels with C$^{18}$O integrated intensity larger than three times the integrated noise are used in the R$_{13,18}$ calculation. Totally 161 pixels in group 1 clouds and 46 pixels in group 2 clouds satisfy this detection criterion. The R$_{13,18}$ values for group 1 clouds lie in the range 7.7-40.2 with a mean of 14.3. For group 2 clouds, the values lie in the range 5.0-21.0 with a mean of 12.0. No significant difference exists between the two groups in the R$_{13,18}$ value. We can see that clouds in both groups possess $^{13}$CO to C$^{18}$O abundance ratios higher than the solar system value (5.5) by a factor of 2.2-2.6, showing that the UV photons from the NGC 2244 and NGC 2237 OB clusters have strong influence on the chemistry of the clouds in the region, although clouds in group 1, as discussed in Section 4.1, exhibit no dynamical effects from the nearby OB clusters. We note that our estimates of $^{13}$CO to C$^{18}$O abundance ratio for the RMC complex are similar to the results by \citep{2014A&A...564A..68S} for the photon-dominated regions (PDRs) in the Orion A giant molecular clouds where they find values lying in the range of 5.7-33.0 with a mean of 16.5. From Figure \ref{fig19} we can see that there is no trend for the abundance ratio with the projection distance from the NGC 2244 OB cluster. This may reflect that the projection distance does not well represent the true distance. We note that all pixels with the $^{13}$CO to C$^{18}$O abundance ratio larger than 25 belong to the group 1 clouds and they are all located around the embedded clusters REFL9 and PL5. We attribute the high isotopologue ratios in that region to the presence of embedded clusters. | 18 | 8 | 1808.09189 |
1808 | 1808.09140_arXiv.txt | Diffuse radio emission from galaxy clusters in the form of radio halos and relics are tracers of the shocks and turbulence in the intra-cluster medium. The imprints of the physical processes that govern their origin and evolution can be found in their radio morphologies and spectra. The role of mildly relativistic population of electrons may be crucial for the acceleration mechanisms to work efficiently. Low frequency observations with telescopes that allow imaging of extended sources over a broad range of low frequencies ($< 2$ GHz) offer the best tools to study these sources. I will review the Giant Metrewave Radio Telescope (GMRT) observations in the past few years that have led to: i) statistical studies of large samples of galaxy clusters, ii) opening of the discovery space in low mass clusters and iii) tracing the spectra of seed relativistic electrons using the Upgraded GMRT. | \begin{figure}[t] \begin{center} \includegraphics[trim= 0.0cm 0.0cm 0.0cm 0.0cm,clip,height=2.15in]{ds9-p200-xray-235.pdf} \includegraphics[trim= 0.0cm 0.0cm 0.0cm 0.0cm,clip,height=2.15in]{ds9-a3376-iau-proc.pdf} \\ \includegraphics[trim= 0.0cm 0.0cm 0.0cm 0.0cm,clip,height=2.0in]{plck171-jvla.pdf} \includegraphics[trim= 0.0cm 0.0cm 0.0cm 0.0cm,clip,height=2.0in]{Venn.pdf} \caption{The single radio relic in the cluster PLCK G200.9-28.2 (top left), the double radio relic in Abell 3376 (top right) and the radio halo in PLCKG171.9-40.7 (bottom left) are shown in white radio contours overlaid on the respective X-ray images of the host clusters in colour. The relics are elongated arc-like sources at the edges of the clusters and a radio halo is a centrally located Mpc-scale extended radio source. Bottom right:- The Venn diagram represents the currently known clusters that are host to one or more kinds of these emission. There are significant number of clusters that are host to one or more kinds of such radio emission. The sample is taken from \citet{2015ApJ...813...77Y}.} \label{eg} \end{center} \end{figure} Clusters of galaxies are gravitationally bound systems of masses $\sim 10^{14}-10^{15}$ M$_{\odot}$ that are composed of dark matter, galaxies and the intra-cluster medium (ICM). The ICM is the most massive baryonic component forming $10-15\%$ of the total mass and mainly contains thermal plasma that emits in X-rays via thermal Bremsstrahlung mechanism. It also contains non-thermal components such as the magnetic fields and relativistic particles but these elude detection in most observing bands. The detection of diffuse radio emission of synchrotron origin from the ICM provides direct evidence for the presence of relativistic electrons ($\sim $GeV) and magnetic fields ($\sim 0.1 - $ a few $\mu$G) in galaxy clusters. These sources are typically classified into radio halos and radio relics based on their morphology and location relative to the X-ray emitting thermal ICM \citep[see][for reviews]{fer12, bru14}. Due to their steep spectra ($\alpha > 1.0$, $\rm S_{\nu} \propto \nu^{-\alpha}$) these sources are typically studied at low radio frequencies ($\leq 2 $GHz). The short radiative lifetime ($\sim 0.1$ Gyr) and long diffusion times ($\geq $Gyr to reach Mpc distance) of relativistic electrons in the ICM requires that radio halos and relics have mechanisms of in-situ re-acceleration associated with their origin \citep[e. g.][]{jaf77}. Radio relics are elongated or arc-like, polarized radio sources that are found at the peripheries of clusters. These occur as single or sometimes in pairs around galaxy clusters (Fig.~\ref{eg}, top). Radio relics are proposed to trace shocks at the cluster outskirts where particles are accelerated \citep[e.g.][]{ens98}. The polarization indicating aligned magnetic fields, spectral indices showing steepening from outer to the inner edges of the relics \citep[e. g.][]{gia08,bon09,wee10,kal12} and the co-spatiality with X-ray detected shocks \citep[e. g.][]{aka12,ogr13} provide support for relics as tracers of shocks. Radio halos are centrally located in the cluster, Mpc sized and unpolarized (Fig.\ref{eg}, bottom left). The origin of radio halos had been proposed to be in the secondary electrons generated by the hadronic collisions in the ICM \citep[e. g.][]{den80,bla99}. The stringent upper limits on the gamma rays associated with this process \citep[e. g.][]{ack10,arlen12} have led to a scenario where a primary mechanism such as turbulent reacceleration \citep[e. g.][]{pet01,bru01} may be playing a crucial role \citep[e. g.][]{bru17}. An empirical scaling relation between the cluster mass (or X-ray luminosity) and the radio power of radio halos is known \citep[e. g.][]{cas13}. The connection between cluster mergers and occurrence of radio halos and relics has been found observationally \citep{buo01,cas10,kal15}. Cluster mergers are a natural origin for the shocks and turbulence that are proposed to play a role in the generation of such sources. Indeed a significant fraction of clusters that host radio halos and relics show presence of both types of sources (Fig.~\ref{eg}, bottom right). However it is still a matter of investigation as to why some merging clusters host radio halos and relics while others do not. For the reacceleration mechanisms to work, a seed population of relativisitic electrons is needed as the efficiencies of acceleration are low \citep[e. g.][]{mar05,kan11,pin17}. These seeds may be due to the radio galaxies in clusters and the secondary electrons. Low frequency observations are crucial in order to trace the seed population as it is expected to be aged synchrotron plasma with steep spectra. In this review I will focus on how low frequency observations have helped to get insights into the occurrence statistics of radio halos and relics, led to the discoveries of such diffuse sources in low mass clusters and are helping to trace the seed populations of relativistic electrons. | Radio halos and relics are direct probes of the non-thermal components in the ICM. Sensitive low frequency observations possible with the current and upcoming interferometers are well suited to reveal the spectra and morphologies of these sources to constrain the theoretical models. In recent years the role of reacceleration by turbulence and shocks in the formation of radio halos and relics has gained considerable support. The availability of seed relativistic electrons so that the reacceleration mechanisms are efficient may be one of the crucial factor dividing merging clusters with and without radio halos and relics. The uniform surveys of galaxy clusters, namely the EGRHS and the ongoing survey mass-complete sample are providing the occurrence statistics of radio halos and relics. The unexplored regime of low mass merging clusters is becoming accessible due to the sensitive low frequency telescopes such as the MWA (see M. Johnston-Hollitt, this volume). The seed relativistic electron population is likely going to be steep spectrum and thus their characterisation relies on sensitive measurements in sub-GHz frequency ranges. The uGMRT is operational and we have carried out the first study of the spectrum of a remnant radio galaxy in the cluster A4038. The prospects to study radio halos and relics with the uGMRT are promising. | 18 | 8 | 1808.09140 |
1808 | 1808.03581_arXiv.txt | A major point of interest in cometary plasma physics has been the diamagnetic cavity, an unmagnetised region in the inner-most part of the coma. Here, we combine Langmuir and Mutual Impedance Probe measurements to investigate ion velocities and electron temperatures in the diamagnetic cavity of comet 67P, probed by the Rosetta spacecraft. We find ion velocities generally in the range 2-4 km/s, significantly above the expected neutral velocity $\lesssim$1~km/s, showing that the ions are (partially) decoupled from the neutrals, indicating that ion-neutral drag was not responsible for balancing the outside magnetic pressure. Observations of clear wake effects on one of the Langmuir probes showed that the ion flow was close to radial and supersonic, at least w.r.t. the perpendicular temperature, inside the cavity and possibly in the surrounding region as well. We observed spacecraft potentials $\lesssim$-5~V throughout the cavity, showing that a population of warm ($\sim$5~eV) electrons was present throughout the parts of the cavity reached by Rosetta. Also, a population of cold ($\lesssim0.1$~eV) electrons was consistently observed throughout the cavity, but less consistently in the surrounding region, suggesting that while Rosetta never entered a region of collisionally coupled electrons, such a region was possibly not far away during the cavity crossings. | \label{sec:intro} \subsection{The Rosetta mission} Between August 2014 and September 2016, the European Space Agency's Rosetta spacecraft followed the short-period, Jupiter Family comet 67P/Churyumov-Gerasimenko in its orbit around the Sun \citep{Glassmeier2007,Taylor2017}. The comet heliocentric distance ranged from 3.5 au at arrival of the spacecraft, to 1.26 au at perihelion in September 2015, to 3.83 au at the end of mission (EOM). The cometocentric distance of the spacecraft was typically on the order of a few tens to a few hundreds of kilometers, providing unprecedented access to the inner coma of a comet. The relative speed of the spacecraft w.r.t.\ the nucleus was generally on the order of one meter per second or less. Previous cometary space missions (e.g.\ ICE at 21P/Giacobini-Zinner \citep{Rosenvinge1986}, Giotto, Sakigake/Suisei and VEGA (1\&2) at 1P/Halley \citep{Reinhard1986,Hirao1987,Sagdeev1987} and Giotto at 26P/Grigg-Skjellerup \citep{Grensemann1993}, which all carried plasma instruments) have been short flybys at distances of at least a few hundred km (and relative speeds of tens of kilometers per second). Thus, the Rosetta mission was unprecedented also with regard to its prolonged stay at the target comet, for the first time allowing the long-term evolution of a comet to be observed by in-situ measurements. \subsection{The diamagnetic cavity} A major point of interest in cometary plasma physics has been the existence, extent and formation mechanism of the diamagnetic cavity, a region in the inner-most part of the coma into which the interplanetary magnetic field cannot reach and which, in the absence of an intrinsic magnetic field of the nucleus \citep{Auster2015}, will be magnetic-field-free. First predicted theoretically by \citet{Biermann1967}, it has since been observed in situ by the Giotto spacecraft at comet 1P/Halley \citep{Neubauer1986}. \citeauthor{Biermann1967} proposed a pressure balance between the magnetic pressure on the outside of the cavity and the ion dynamic pressure on the inside to account for its formation and extent. However, \citet{Ip1987} and \citet{Cravens1986,Cravens1987} found this to be insufficient to explain the extent of the cavity observed at Halley and instead invoked the ion-neutral drag force inside the cavity to balance the outside magnetic pressure. This was supported by observations of near-equal ion and neutral velocities inside the cavity ($\sim$1~km/s and $\sim$0.9~km/s, respectively), consistent with strong ion-neutral collisional coupling, and clear stagnation of the ion flow in the region just outside the cavity boundary \citep{Balsiger1986,Krankowsky1986}. A diamagnetic cavity was first detected around comet 67P in magnetometer (RPC-MAG \citep{Glassmeier2007}, hereafter MAG) data from July 26, 2015 \citep{Goetz2016a}, near perihelion at a distance of 170 km from the nucleus (in the terminator plane). The spacecraft remained inside the cavity for about 25 min during this event. Subsequent analysis has identified a total of 665 cavity crossings in MAG data \citep{Goetz2016b}, between April 2015 and February 2016 (i.e.\ some preceding the original detection). They ranged in duration from 8 s up to 40 min, in distance to the nucleus from 40 to 380 km and in heliocentric distance from 1.25-2.4 au. The low velocity of Rosetta ($\lesssim 1$ m/s) implies that these highly transient events were the result of the cavity expanding and contracting over Rosetta's position, rather than resulting from the spacecraft moving into and out of a stationary cavity. Another possibility is blobs of unmagnetized plasma detaching from the main cavity structure and convecting past the spacecraft. The distance to the nucleus of the cavity crossings exhibited a strong statistical dependence on the long-term production rate, but was unaffected by diurnal variations and short-duration events such as outbursts or varying solar wind consitions. \citet{Goetz2016a} therefore suggested a Kelvin-Helmholtz type instability, driven by a presumed velocity shear at the cavity boundary, to account for its short-term dynamics. This was also proposed to explain the fact that cavity distances were generally found to be larger than predicted for a steady-state cavity sustained by the above pressure balance, as in hybrid simulations by \citet{Koenders2015} and \citet{Rubin2012}. The existence of instabilities at the cavity boundary was indeed confirmed in these simulations. Density measurements by the Mutual Impedance Probe (RPC-MIP \citep{Trotignon2007}, hereafter MIP) inside the diamagnetic cavity showed densities ranging from $\sim$100 to $\sim$1500 cm$^{-3}$ on longer time scales, but that were almost constant inside any given cavity or between closely successive events \citep{Henri2017}. The surrounding regions of magnetized plasma were in contrast characterized by large density variations, predominantly in the form of large-amplitude compressible structures with relative density fluctuations $\delta n/n \sim$ 1 \citep{Harja2018}. These generally matched similar structures observed in the magnetic field near the cavity by \citet{Goetz2016a}. The plasma density inside the cavity was found to be entirely determined by the ionization of the cometary neutral atmosphere and the cavity boundary generally located close to the electron exobase. Hence, \citeauthor{Henri2017} suggested that the cavity formation and extent was the result of electron-neutral collisionality rather than the ion-neutral collisionality previously invoked. They also proposed a Rayleigh-Taylor type instability of the cavity boundary, driven by the electron-neutral drag force acting as an "effective gravity", instead of the Kelvin-Helmholtz type suggested by \citet{Goetz2016a}. % \citet{Timar2017} obtained good fits of observed cavity distance values to the ion-neutral drag model of \citet{Cravens1986,Cravens1987}, using the cometary neutral production rate and solar wind dynamic pressure estimated from magnetic field data as well as several different solar wind propagation models. This possibly eliminates the need for an instability at the cavity boundary to account for the intermittent nature of the cavity crossings, in favour of variations in the solar wind pressure. \citet{Nemeth2016} found that accelerated electrons in the 100~eV range, typically present in the inner coma, were absent inside the diamagnetic cavity, suggesting that these electrons were bound to the field lines and therefore excluded from the cavity. \subsection{Neutral gas velocity} For the first few months at the comet, through fall of 2014 up to early 2015, empirical or semi-empirical estimates of the expansion velocity of the neutral coma gas are available from many different sources: doppler shift of the spectral lines of water observed by the Microwave Instrument on the Rosetta Orbiter (MIRO, \citet{Gulkis2007}) \citep{Lee2015,Biver2015,Gulkis2015}, simulation outputs of DSMC models \citep{Bieler2015} constrained by ROSINA-DFMS data \citep{Fougere2016a,Fougere2016b} and direct measurements by ROSINA-COPS ram and nude gauges \citep{Tzou2017}. They all typically give terminal velocities at a few kilometers from the nucleus surface in the range 400~-~800~m/s, with generally a positive correlation between velocity and local outgassing intensity. However, from the period between April 2015 and February 2016 considered in this Paper, published measurements are scarce. \citet{Marshall2017} used a range of 400~m/s to 1~km/s, with a preferred value of 700~m/s, to obtain local effective production rates of H$_2$O from H$_2^{16}$O/H$_2^{18}$O line area ratios obtained by MIRO for the entire period from August 2014 to April 2016, though no more specific values are given from within this period. \citet{Heritier2017} used a one-dimensional model for the neutral gas based on an adiabatic fluid expansion around the nucleus driven by inner boundary conditions on gas outflow velocity from \citet{Huebner2000} and temperature from the thermophysical model of \citet{Davidsson2005} to find terminal velocities of about 800~m/s. For the purpose of comparison to observed ion velocities in this study, we note simply that the neutral outgassing velocity is on the order of 1~km/s, and that exact values are likely to be lower than this estimate rather than higher. \subsection{Ion velocity} The primary ionization processes in the cometary coma, photoionization and electron impact ionization \citep{Vigren&Galand2013,Galand2016}, produce ions that are initially cold and flowing with the neutral gas. This is an effect of conservation of momentum: the momentum of the ionizing particle is minuscule compared to that of the much heavier neutral molecule and therefore does not affect its motion in any significant way. The excess energy from the ionization instead goes to the electrons, which are therefore born warm ($T_{\textnormal{e}}$ $\sim$ 10 eV) \citep{Häberli1996,Galand2016}. If there is no electric field, or if the ions are strongly collisionally coupled to the neutrals, the ions can thus be expected to be cold and flowing with the neutral gas. In the presence of a magnetic field of solar wind origin, the assumption of no electric field fails because of the existence of a convective electric field, which will cause the ions to gyrate, $\mathbf{E} \times \mathbf{B}$ drift and eventually become dynamically part of the solar wind flow (which will be decelerated and deflected by mass-loading) \citep{Coates2004,Szego2000}. This so called \emph{ion pick-up} process takes place over spatial scales on the order of the ion gyro-radius, which for singly charged water group ions ($m_{\textnormal{i}} \approx 18$) in a cometary plasma with typical magnetic field strength $\lesssim$20 nT \citep{Goetz2017} is $\gtrsim$10 km for ion velocities $\gtrsim$1 km/s. Inside the diamagnetic cavity, or at distances outside of it smaller than about 10 km, this process is therefore unimportant for the ion motion. However, the presence of warm electrons (to be further discussed below) suggests the existence of an ambipolar electric field (at least inside the diamagnetic cavity) of a strength on the order of $k_{\textnormal{B}} T_{\textnormal{e}}/q_{\textnormal{e}}r$ to maintain quasi-neutrality of the radially expanding cometary plasma. In such a case, strong collisional coupling to the neutrals is necessary if the ions are to remain at the neutral velocity. Estimates of the location of the ion-neutral collisionopause by \citet{Mandt2016} suggested that Rosetta was generally in a region where ion-neutral collisions were important. However, these estimates did not take into account the reduced collisionality of accelerated ions due to the cross-section for ion-neutral collisions decreasing with energy. \citet{Vigren2017} used a 1D model to simulate the radial acceleration of water group ions interrupted by collisions (primarily charge transfer processes) with neutral water molecules, taking into account the energy-dependence of the cross-sections. They found that for an outgassing rate $\sim$2$\cdot 10^{28}$ s$^{-1}$, typical of 67P near perihelion, even a weak electric field of 0.03 mV/m, typical of what would be expected for an ambipolar field, is sufficient to partially decouple the ions from the neutrals, giving a bulk ion velocity of about 4 km/s at distances $\sim$200 km from the nucleus, typical of the Rosetta spacecraft around perihelion. For an outgassing rate $\sim$2$\cdot 10^{29}$ s$^{-1}$, typical of Halley during the Giotto encounter, collisional coupling was found to prevail. \citet{Vigren2017b} combined Langmuir probe (RPC-LAP \citep{Eriksson2007}, hereafter LAP) and MIP measurements to produce estimates of the ion velocity for a three-day period near perihelion, including one diamagnetic cavity crossing. They obtained values typically in the range 2-8~km/s, roughly in line with the predictions of \citet{Vigren2017}, lending further support to the supposition that ions are collisionally decoupled from the neutrals at 67P. The presence of an ambipolar electric field, the velocity of the ions and the formation and dynamics of the diamagnetic cavity at 67P are at present poorly understood. In this Paper we attempt to shed some light on these issues by using the method of combined LAP and MIP measurements to produce estimates of the ion velocity throughout the diamagnetic cavity and compare to the surrounding region. \subsection{Electron temperature} \citet{Odelstad2015,Odelstad2017} presented measurements of the spacecraft potential ($V_{\textnormal{S/C}}$) by LAP, showing that $V_{\textnormal{S/C}}$ was mostly negative throughout Rosetta's stay at the comet, often below -10 V and sometimes below -20 V. This was attributed to a population of warm ($\sim$5-10 eV) coma photoelectrons, whose presence was explained by the neutral gas not being dense enough to effectively cool these electrons (which are born warm, as mentioned above) by collisions. Positive spacecraft potentials ($\sim$0-5 V) were only observed in regions of very low electron density ($\lesssim$10 cm$^{-3}$), typically far from the nucleus or above the more inactive areas on it, where the positive $V_{\textnormal{S/C}}$ could be explained by low density rather than temperature and where significant electron cooling by neutrals was not possible. Thus, it was concluded that such warm electrons were persistently present in the parts of the coma reached by Rosetta, most notably also around perihelion, where the elevated neutral density would perhaps have been expected to effectively cool the electrons. The statistical nature of this study could not rule out the existence of some brief events of low spacecraft potential hiding in the data set, which would indicate the near-absence of warm electrons. In this paper, we examine this in detail for the diamagnetic cavity crossings and discuss the implications for the physics of the cavity. In addition to these warm electrons, clear signatures of cold ($\lesssim$0.1 eV) electrons have also been observed by LAP \citep{Eriksson2017} and MIP \citep{Gilet2017}. In LAP, these show up in high-time-resolution current measurements at fixed bias voltage in the form pulses of typical duration between a few seconds and a few minutes, and in bias voltage sweeps in the form of very steep slopes in the current-voltage curve at high positive bias voltages (to be discussed further below). In MIP, they produce a second resonance in the mutual impedance spectra below the total plasma frequency. Since local electron cooling was negligible as evidenced by the presence of warm electrons, this cold plasma was inferred to have formed in a region closer to the nucleus than reached by Rosetta. Together with the intermittent nature of the signatures in LAP data, this was taken as evidence for strong filamentation of the cold plasma close to the nucleus, with individual filaments extending far outside the collisionally dominated region, perhaps even detaching from it entirely. Similar structures were observed to develop in connection to the diamagnetic cavity in global 3D hybrid simulations by \cite{Koenders2015} and have been proposed to result from the instability of the cavity boundary (e.g.\ \citet{Henri2017}). However, observations of cold electrons are not limited to the diamagnetic cavity so the relationship (if any) between the structure and dynamics of cold and unmagnetized plasma, respectively, is not clear. In this paper, we investigate this by examining in detail the presence of cold electrons in the diamagnetic cavity and the surrounding region. | \label{sec:conclusions} We have combined density measurements by MIP with ion slopes from LAP sweeps to produce measurements of the ion velocity in and around the diamagnetic cavity of 67P. As noted before by \citet{Henri2017}, the density inside the cavity was remarkably stable, while the surrounding region was characterized by high density variations. This turned out to hold also for the ion slopes; they were much more variable outside than inside the cavity. The resulting ion velocity measurements inherited this feature, with stable velocities narrowly distributed around 3.5-4 km/s inside the cavity and scattered toward higher velocities ($\lesssim$8-10 km/s, though still heavily weighted towards 4 km/s) in the surrounding region. The ion velocity inside the cavity was clearly elevated w.r.t.\ the neutral velocity of $\lesssim$1 km/s and in good agreement with the model predictions of \citet{Vigren2017}. This indicates that the ions were not collisionally coupled to the neutrals and implies that the ion-neutral drag force was not responsible for balancing the outside magnetic pressure at the cavity boundary. It also suggests the existence of an ambipolar electric field to accelerate the ions, at least inside the cavity. Clear evidence of wake effects on LAP2 has been identified inside the cavity during a spacecraft slew. The geometry of this slew was such that LAP2 was brought forward from a position close to being obscured from view of the nucleus to a position more exposed to radial flow from the comet. Inside the cavity, the effect of this on LAP2 was to increase its collection of ions and cold electrons and generally make its current-voltage characteristics more similar to that of LAP1, which was consistently well-exposed to any radial flows. This shows that the flow of cometary ions inside the cavity was indeed close to radial and supersonic, at least w.r.t.\ the perpendicular temperature. Thus, the observed ion velocities inside the cavity are to be taken as a radial velocity of the ions, possibly a cold radial drift, although a non-negligible radial temperature of the ions cannot be rule out. Outside the cavity, the effects of these slews on the LAP2 current were less clear and consistent, indicative of a more dynamic wake, or possibly a more variable ion flow direction. We also made a detailed examination of the spacecraft potential measurements inside the cavity, finding it to be consistently $\lesssim$-5 V. This proves that a population of warm ($\sim$5-10 eV) electrons was present throughout the parts of the cavity probed by Rosetta. This was shown to hold on larger time scales throughout the mission already by \citet{Odelstad2017}, but has now been confirmed on a sweep-by-sweep level for all cavity events. This shows that Rosetta never entered the region of collisionally coupled electrons presumed to exist in the innermost part of the coma, not even during any of the passes through the diamagnetic cavity. Finally, we have found that a population of cold ($\lesssim$ 0.1 eV) electrons, first shown by \citet{Eriksson2017} to be intermittently present at Rosetta, is in fact observed consistently throughout the diamagnetic cavity. Already immediately outside the cavity, sweeps lacking a signature of cold electrons begin to turn up intermittently. This reinforces the notion of \citet{Henri2017} that the formation and extent of the cavity is closely related to electron-neutral collisionality and suggests that the the region of collisionally coupled electrons, while never entered by Rosetta, was possibly not that far away during the cavity crossings. It also suggests that the filamentation of the cold electrons begins at or near the cavity boundary and is possibly related to an instability of this boundary. \appendix | 18 | 8 | 1808.03581 |
1808 | 1808.03853_arXiv.txt | We present the modeling tool we developed to incorporate multi-technique observations of Cepheids in a single pulsation model: the Spectro-Photo-Interferometry of Pulsating Stars (SPIPS). The combination of angular diameters from optical interferometry, radial velocities and photometry with the coming Gaia DR2 parallaxes of nearby Galactic Cepheids will soon enable us to calibrate the projection factor of the classical Parallax-of-Pulsation method. This will extend its applicability to Cepheids too distant for accurate Gaia parallax measurements, and allow us to precisely calibrate the Leavitt law's zero point. As an example application, we present the SPIPS model of the long-period Cepheid RS Pup that provides a measurement of its projection factor, using the independent distance estimated from its light echoes. | One century after the discovery of their Period-Luminosity relation (the Leavitt law) by \cite[Leavitt \& Pickering~(1912)]{Leavitt12}, Cepheids are still the keystone of the empirical cosmic distance ladder. However, due to the large distances of Cepheids (particularly the long-period pulsators), only relatively imprecise parallax measurements are available (e.g., from Hipparcos; \cite[van Leeuwen \etal\ ~2007]{vanLeeuwen07}). As a result, the calibration of the zero point of the Leavitt law using Galactic Cepheids is insufficient accurate. Gaia's DR2 high accuracy parallaxes will bring a tremendous improvement, with distances to hundreds of individual Cepheids measured to a few percent accuracy. Alternatively, a classical avenue to measure the distances to individual Galactic and LMC Cepheids is the Parallax-of-Pulsation (PoP) method, also known as the Baade-Wesselink (BW) technique. We here briefly present the SPIPS modeling tool that we developed to reproduce the classical observables of the pulsation of Cepheids (radial velocity, photometry, interferometry,...), and how we plan to use it in conjunction with the Gaia parallaxes to calibrate the PoP technique and eventually the Leavitt law. | 18 | 8 | 1808.03853 |
|
1808 | 1808.08236_arXiv.txt | {We demonstrate the existence of a generic, efficient and purely gravitational channel producing a significant abundance of dark relics during reheating after the end of inflation. The mechanism is present for any inert scalar with the non-minimal curvature coupling $\xi R\chi^2$ and the relic production is efficient for modest values $\xi = {\cal O}(1)$. The observed dark matter abundance can be reached for a broad range of relic masses extending from $m \sim 1 {\rm k eV}$ to $m \sim 10^{8} {\rm GeV}$, depending on the scale of inflation and the dark sector couplings. Frustratingly, such relics escape direct, indirect and collider searches since no non-gravitational couplings to visible matter are needed. } \emailAdd{[email protected]} \emailAdd{[email protected]} \emailAdd{[email protected]} \emailAdd{[email protected]} \begin{document} \begin{flushleft} \hfill IMPERIAL/TP/2018/TM/04\\ \hfill KCL/PH-TH-2018-56 \end{flushleft} | Dark matter may consist of particles which were never in chemical or kinetic equilibrium with visible matter, in contrast to thermal relics \cite{Bertone:2004pz}. The coupling of the dark sector to visible matter may be too weak to maintain equilibrium but still large enough to generate the relic abundance through out of equilibrium decays of the visible matter \cite{McDonald:2001vt}. This is commonly referred to as the freeze-in mechanism, or FIMP dark matter. See \cite{Bernal:2017kxu} for a recent review. Dark matter may also be completely decoupled from the visible matter and interact only gravitationally. A well known example is the WIMPZILLA scenario \cite{Kolb:1998ki, Chung:2001cb} where dark matter particles are produced gravitationally at the end of inflation\footnote{In \cite{Kolb:1998ki} WIMPZILLA refers to any non-thermal superheavy dark matter particle produced either gravitationally or by direct couplings, such as inflaton decays during preheating and reheating \cite{Kolb:1998ki, Chung:1998bt}.} and must be superheavy to yield the observed relic abundance. Perturbative gravitational production may also proceed through graviton mediated scatterings after the end of inflation \cite{Garny:2015sjg, Tang:2016vch}. Moreover inflationary fuctuations of light spectators scalars, completely decoupled from the visible matter, may contribute to dark matter. This however generates isocurvature dark matter \cite{Linde:1985yf,Nurmi:2015ema} heavily constrained by observations \cite{Ade:2015lrj}. Yet another, efficient and purely gravitational channel producing adiabatic dark matter was recently discovered in \cite{Markkanen:2015xuw}. Subsequently similar setups were further explored in \cite{Ema:2016hlw}. This relies on the non-minimal coupling $\xi R \chi^2$ which, under very generic conditions, is generated by radiative corrections for any energetically subdominant spectator scalar $\chi$ \cite{Chernikov:1968zm}. During reheating, the universe is dominated by the oscillating inflaton field and the scalar curvature $R$ is an oscillatory function which periodically takes negative values. When $R$ is negative the $\chi$ particles have negative mass squared. This results in an instability and a very efficient particle production, similarly to the cases of tachyonic preheating \cite{Bassett:1997az,Tsujikawa:1999jh,Dufaux:2006ee,Bassett:1999mt, Finelli:1998bu} and vacuum instability \cite{Herranen:2015ima,Kohri:2016wof,Postma:2017hbk}. If the produced $\chi$ particles are stable and decoupled from visible matter they may constitute a natural dark matter component. We call these {\it despicable dark relics} for two reasons: the mechanism is very generic and in the absence of non-gravitational interactions the relics would escape all direct, indirect and collider searches of dark matter\footnote{We do not consider the effects due to possible gravitational breaking of global symmetries~\cite{Kallosh:1995hi}.}. The particle production is efficient already for modest values of the non-minimal coupling $\xi = {\cal O}(1)$ and the relic mass window spans several orders of magnitude extending down to sub-keV scales \cite{Markkanen:2015xuw}. This is a significant difference compared to gravitationally produced WIMPZILLAs \cite{Kolb:1998ki, Chung:2001cb} which must be superheavy $m \gtrsim 10^{12}$ GeV to yield the observed dark matter abundance\footnote{Note that, like our setup, the graviton mediated scatterings discussed in \cite{Garny:2015sjg, Tang:2016vch} can also produce a significant abundance of light relics. However, this requires a very efficient reheating whereas our setup is insensitive for the duration of reheating. It is possible that both mechanisms contribute simultaneously to the dark matter.}. If there are no isocurvature perturbations at the end of inflaton, all regions in the observable universe go through the same reheating history and acquire the same abundance of despicable relics \cite{Markkanen:2015xuw}. Therefore, the dark matter generated through the non-minimal coupling is adiabatic as required by observations \cite{Ade:2015lrj}. In this work we perform a detailed investigation of the despicable dark relics. We assume the dark sector to consist of a non-minimally coupled scalar which may have a self-interaction $\lambda \chi^4$ but no couplings to visible matter. We consider three different types of reheating stage characterised by an equation of state $w(t)$ which corresponds to inflaton oscillations in a quadratic and a quartic potential, and a kination dominated epoch. In each case we compute the final dark matter yield and find that the observed dark matter abundance can be easily obtained. In fact, the mechanism is so efficient that the spectator couplings are subject to non-trivial constraints to avoid overproduction of dark matter. The paper is organised as follows. In Section \ref{sec:tach} we review the analytical formalism for tachyonic particle production and in Section \ref{sec:num} we solve the equations numerically. In Section \ref{sec:ev} we study the evolution of particle number in the presence of dark matter self-interactions and present our main results for the final dark matter yield. Finally, in Section \ref{sec:res} we present our conclusions. Our sign choices are (+,+,+) in the classification of \cite{Misner:1974qy}. | In this work we have explored a mechanism, first set out in \cite{Markkanen:2015xuw}, where dark matter is produced gravitationally during reheating. We find that the mechanism is very efficient, and the observed dark matter abundance can be reached for a very broad range of relic masses. We have investigated three different reheating scenarios, corresponding to inflaton oscillations in quadratic and quartic potentials, and a kination epoch. In all the cases the curvature scalar $R$ evolves to negative values which may trigger a tachyonic instability for scalars with the non-minimal coupling $\xi R \chi^2$. If the produced particles are stable, they constitute adiabatic dark matter \cite{Markkanen:2015xuw}. We have focused on a dark sector consisting of a single non-minimally coupled scalar $\chi$ with a mass $m$ and a self-interaction $\lambda \chi^4$, and no couplings to any other matter fields. We have concentrated on the region $\xi \gtrsim 1$ where the scalar does not fluctuate during inflation but can experience a strong tachyonic instability during reheating. We have performed a detailed numerical analysis of the particle production for each scenario, varying the reheating scale, its duration, the non-minimal coupling $\xi$ and the self-interaction strength $\lambda$ which all affect the yield. We have followed the evolution of the relic particle number from the end of reheating until today, accounting in particular for the possible impacts of inelastic scatterings mediated through the self-coupling $\lambda \chi^4$. The main results of this work are encompassed in Figures \ref{fig:quadraticlam}, \ref{fig:quarticlam} and \ref{fig:kinationlam} which show the present day abundance of gravitational relics. The observed dark matter abundance can be obtained in broad range of relic masses, extending down to keV scale, and for realistic dark sector couplings $\lambda$ and $\xi$. This should be contrasted to gravitationally produced WIMP-ZILLAs which must be superheavy $m\gtrsim 10^{12}$ GeV to yield the observed abundance \cite{Kolb:1998ki, Chung:2001cb}. The difference stems from the efficiency of the tachyonic particle production mechanism. For $\xi\gg 1$ the mechanism can in fact easily yield too much dark matter, which opens up a potentially interesting new way to constrain theories with stable extra scalars. The key features of the scenario are its genericity and that the relics can be completely decoupled from visible matter. The fact that they would thus remain undetectable in all conceivable dark matter searches, motivates us to introduce the name Despicable Dark Relics (DDR) for this type of dark matter. | 18 | 8 | 1808.08236 |
1808 | 1808.00468_arXiv.txt | Effective stellar feedback is used in models of galaxy formation to drive realistic galaxy evolution. Models typically include energy injection from supernovae as the dominant form of stellar feedback, often in some form of sub-grid recipe. However, it has been recently suggested that pre-SN feedback (stellar winds or radiation) is necessary in high-resolution simulations of galaxy evolution to properly regulate star formation and properties of the interstellar medium (ISM). Following these processes is computationally challenging, so many prescriptions model this feedback approximately, accounting for the local destruction of dense gas clouds around newly formed stars in lieu of a full radiative transfer calculation. In this work we examine high resolution simulations (1.8~pc) of an isolated dwarf galaxy with detailed stellar feedback tracked on a star-by-star basis. By following stellar ionizing radiation with an adaptive ray-tracing radiative transfer method, we test its importance in regulating star formation and driving outflows in this galaxy. We find that including ionizing radiation reduces the star formation rate (SFR) by over a factor of 5, and is necessary to produce the ISM conditions needed for supernovae to drive significant outflows. We find that a localized approximation for radiation feedback is sufficient to regulate the SFR on short timescales, but does not allow significant outflows. Short and long range radiation effects are both important in driving the evolution of our low-metallicity, low-mass dwarf galaxy. Generalizing these results to more massive galaxies would be a valuable avenue of future research. | \label{sec:intro} Historically, simulations of galaxy formation have suffered from the ``overcooling'' problem, whereby the action of self-gravity and radiative cooling alone produces galaxies with far too many stars. This problem has been addressed by employing various models of strong stellar feedback physics which are capable of generating self-regulating star formation in galaxies (see \cite{SomervilleDave2015} and \cite{NaabOstriker2017} for recent reviews). Energy injection from supernovae (SNe) has historically been used as the sole form of feedback. However, this is generally done heuristically as many simulations lack the ability to resolve the Sedov phase of individual SNe. But even with this strong feedback, recent work has argued for the need for pre-SN feedback, from stellar winds and/or stellar radiation \citep[e.g.][]{Hu2016,Hopkins2018}, though this is typically modeled simply as additional energy injection around newly formed stars. The need for additional feedback is confirmed by simulations that are capable of fully resolving individual SN \citep[e.g.][]{Peters2017,Smith2018a,Smith2018b,Hu2018}. However, modeling these processes in detail is challenging, and their competing effects on galactic evolution are poorly constrained. Radiation from massive stars dominates the total feedback energy output of a stellar population \citep[e.g.][]{Abbott1982,Leitherer1999,Agertz2013}, surpassing the energy ejection of supernova ($\sim 10^{51}$~erg) by two orders of magnitude. If radiation couples effectively to the interstellar medium (ISM), it can be a substantial source of additional feedback. Simulations including stellar radiation feedback followed through radiative transfer or radiation-hydrodynamics schemes have found it to be effective in regulating star formation and driving galactic winds \citep[e.g.][]{WiseAbel2012,Kim2013a,Sales2014,Oshea2015,Rosdahl2015,Pawlik2015,Ocvirk2016,Peters2017}. This occurs in four ways: 1) heating of diffuse gas and preventing the formation of cold, dense star formation regions, 2) destruction of cold, dense gas around recently formed stars, preventing further star formation, 3) momentum input by direct absorption of UV radiation by gas and (in some cases) dust through re-emission and scattering in the infrared, and 4) lowering the typical ISM densities in which SNe occur and greatly increasing their effectiveness. However, most works that employ stellar radiation feedback to account for these effects do so using various forms of sub-grid, approximate models to avoid the substantial additional cost of full radiative transfer. Many works use a Str{\"o}mgren approximation whereby the particles / cells within the Str{\"o}mgren radius around a radiating star are heated and kept ionized, with additional approximations made to account for overlapping ionized regions \citep[e.g.][]{HQM2011,Hu2016,Hu2017}. Other works employ some form of energy or momentum injection localized to the region immediately around a star particle \citep[e.g.][]{Agertz2013,Roskar2014,Ceverino2014,Forbes2016}. Although some of these approximate methods account for long-range effects of diffuse radiation \citep{HQM2012,Hopkins2018} most cases treat local radiation {\it only}, confined to energy or momentum injection in a limited physical region around a star particle. This is done both because the local destruction of dense clumps of gas around newly formed stars is commonly believed to be the dominant impact of stellar radiation feedback and because it is computationally less expensive to implement. The role of long-range stellar radiation, once ionizing photons break out of star forming regions, is not well characterized. Indeed it remains to be seen if modeling only the short-range effects of stellar radiation feedback is sufficient. In this work we use a detailed model for stellar feedback presented in \cite{Emerick2018} (hereafter Paper I) to study the role of radiation feedback in dwarf galaxy evolution. For the first time in a galaxy-scale simulation, we resolve individual HII regions using an adaptive ray tracing radiative transfer method to follow the ionizing radiation from particles that represent individual stars. We focus on addressing two questions: 1) what role does radiation feedback play in regulating star formation, and 2) are the long-range effects of radiation feedback important, or is the local destruction of dense gas the dominant effect. To investigate these questions, we compare three simulations of the evolution of an isolated, low mass dwarf galaxy. Our fiducial model, containing full stellar radiation feedback, is compared against a run without ionizing radiation feedback, and a run with ionizing feedback limited to the local region around a given star. We discuss our methods in Section~\ref{sec:methods} and present our results ion Section~\ref{sec:results}. | \label{sec:conclusion} In agreement with previous work we find that (local) stellar radiation feedback is effective in regulating star formation, but that non-local ionizing radiation is key for driving outflows in our simulations of an isolated, low mass, dwarf galaxy. Simulations run without ionizing radiation feedback have star formation rates a factor of five higher than our fiducial simulation. Despite the lower rate, SNe in our fiducial run are capable of driving larger galactic outflows, aided significantly by the ionizing radiation feedback. We demonstrate for the first time that simple prescriptions of local stellar radiation feedback fail to reproduce the evolution of our fiducial model. Our simulation with radiation localized to 20~pc around each star particle does effectively regulate star formation on short time scales, predominately by quickly destroying cold, dense gas around young, hot stars. However, this model does not drive the significant outflows seen in our fiducial simulation. Long-range ionizing radiation is important for carving channels allowing the ejection of significant amounts of mass and metals from the SNe. Our simulation with localized radiation feedback retains a significantly higher fraction of metals than expected observationally for low mass dwarf galaxies. Finally, we note that we have performed this experiment on only one possible type of galaxy. Its low virial temperature ($\sim10^{4}$~K) makes this galaxy particularly sensitive to the effects of stellar feedback, and ionizing radiation in particular. Examining the role of long-range, diffuse stellar ionizing radiation on star formation and galactic winds in more massive galaxies is an important avenue of future research. | 18 | 8 | 1808.00468 |
1808 | 1808.10278_arXiv.txt | {The properties of universes are explored that are entirely in the interior of black holes. It is argued that these models offer a paradigm that may shed a new light on old cosmological problems. The topics that are addressed include: geometry of the universes, evolution, relation to the concordance model, dark matter, dark energy, mass ejection from black holes, galaxy models with a central black hole, Mach's principle. } \dedicated{I dedicate this paper to all non-expert people in the world who try to make sense of cosmology, and, with their own means and visions, long to probe its mysteries.} \begin{document} \def\cyc{{\rm cyc}} \def\dacyc{{\rm dacyc}} | It is well-known that a closed universe with a Robertson-Walker metric oscillates from the big bang to a maximum expansion state and back to the big crunch. At the maximum expansion state the radius of the universe equals its Schwarzschild radius $2GM/c^2$, with $M$ the mass of the universe, $G$ the gravitational constant and $c$ the velocity of light.\footnote{See also subsection~\ref{sect_defgen} for the definition of $M$, and subsection~\ref{sect_synchronous}.} This suggests that one may explore universes that are the interiors of black holes. To this idea can be added that tidal forces generated by a spherical mass distribution scale as $M/r^3$, with $r$ the distance to the center of the distribution. This result holds for the Schwarzschild metric as well (\cite{MTW} 1973) in a local geodetic frame (freely infalling observer). Since the Schwarzschild radius scales with $M$, it follows that at that radius the tidal forces for an infalling observer scale as $M^{-2}$. Hence these tidal forces are negligible for a very massive black hole.\footnote{This fact is of no particular importance for this paper, but is mentioned here because that idea germinated the paper. One is used to think of black holes as fierce things that destroy everything that pass the event horizon. Very massive black holes, on the contrary, offer a gentle, but irreversible, welcome.} Of course, at the center of such a thing there waits disaster, but what if, as suggested above, inside the black hole there would be a distributed mass distribution and not a point mass? The density of such a distributed mass distribution with a constant mass density also scales as $M^{-2}$, hence could also be very small. The infalling observer would then simply enter a new part of spacetime, with the only consequence that he/she would never be able to leave it again. Such a thing could be called a universe, and it is the purpose of this paper to show that it actually is a universe much in the same sense as classical cosmology defines it. In this paradigm, it comes natural to consider a 3D universe as embedded in another 3D universe. Since we will show that it can be the interior of a (spherical) black hole, we will adopt spherical symmetry. It is useful to make clear, from the outset, what the geometrical difference is between the universes in this paper and the standard Robertson-Walker (RW) universes. To that end, it helps to reduce the dimension by one, and consider 2D universes. A 2D RW universe is confined to the surface of a sphere, which is embedded in 3-space. In this paper, a 2D universe with positive curvature in all points is the interior of a circle, though the surface inside the circle need not be flat. These universes therefore have a center (a point) and a boundary (a circle). Likewise, the 3D the universes we consider here have a center (a point) and a boundary (the surface of a sphere). As for the center, we can characterize it in an ideal, perfectly ordered universe, which we will call a synchronous universe\footnote{According to this paper's definition of synchronous, the RW models are synchronous.}. It is the only place were a material particle can stay put and not partake in the expansion or contraction. In reality, collisions during the contraction phase will give some momentum to any structure formed, and hence all matter will have velocities with effects that add to the effects of the expansion or contraction.\footnote{The wording here is unusually convoluted, since we will argue that the term 'comoving velocity' is somewhat of a misnomer.} In the same ideal world we will show that the center is the only place where the Hubble parameter is the same in all directions, but we will argue that current observations are not good enough to disentangle local anisotropies from the effects of sphericity, let alone that we would, at present, have a clue where we would be located in such a universe. We will devote considerable attention to the boundary. Suffice here to state that also at the boundary space-time is locally 4D Lorentzian in all directions, as it is everywhere, since we will show that at the boundary the singularity in the radial coefficient of the metric (this is one of the characterizations of the boundary) can be transformed away. In this paper we cannot exclude the now obsolete assumption that the universe could have gone through (many) cyclic phases of expansion and contraction, though we will try to define these contractions carefully and we will not need to assume any particularly neat periodic, nor global, behaviour. During contractions, the space time forced (or facilitated) structure formation, up to the merging of black holes and accretion of matter on them. The resulting dramatic increase in gravitational binding energy, together with increasingly intense background radiation, caused an immense radiation field, which therefore is a necessary companion of the next big bang, or rather, the next--relatively quiet--expansion phase. However, we will not, and need not, assume that a big crunch goes all the way to "the point"; a condensed state, possibly with multiple centres in non-synchronous evolution, suffices. The spherical symmetry (or possible more complex geometries for which spherical symmetry is but the simplest of all models) needs to be reconciled with the isotropy of the Cosmic Microwave Background (CMB). While in standard cosmology inflation is one of the mechanisms to obtain isotropy, acting from the inside out, we will argue that isotropy can also come about by the mixing that the phases of contraction imply, more in particular during the dense phases. This mixing acts, so to speak, from the outside inwards. Hence the universes in this paper are, by design, homogeneously filled with photons (which happen to be microwave photons in the current state of the universe). Of course, some riddles can be present, but unlike in standard cosmology these riddles are probably not particularly informative on the history of the universe. Observations don't seem to indicate that pressure is an important player, and therefore we can suffice with considering pressureless (dusty) universes. This paper is best read by first jumping to the summary. | In this paper, we consider spherically symmetrical cosmological models that are the interior of black holes, for our purposes defined as the matter in the spacetime inside a Schwarzschild horizon. By virtue of the spherical symmetry, the universe can be ordered in spherical shells, which are given a label which is not a radius (e.g. distance to the center, as would be appropriate for an Euclidean 3-sphere), but these labels preserve order with the surface area of the shells. The labels are time independent, in contrast with the surface area (subsection~\ref{sect_defgen}). We employ the well-known Lema\^\i tre-Tolman models (section~\ref{chapt_LT}), but recast them in a form that suits our purposes. More in particular, we introduce in subsection~\ref{sect_definitions} the function 'cyc', which is a generalization of the well known cycloidal solution of the Robertson-Walker (RW) models. We study it extensively in appendix~\ref{sect_cyc}. This function not only packs all technical mathematical details of the time evolution equation into one mathematical object, thereby formally separating all mathematical details from the essence of the physics, but its introduction is also almost necessary in order to deal elegantly with the complications that arise from the spherical symmetry. In subsection~\ref{sect_conn_standard} we briefly show how 'cyc' connects with the standard practices in cosmology and the parameters of the 'concordance' $\Lambda$CDM model. We present the view that, in order to understand the current state of our universe, it is essential to consider states in the very distant past in which the universe was not in expansion. In section~\ref{chapt_early} we muse about these phases, and argue that the isotropy of the Cosmic Microwave Background can be understood as a consequence of them, in a qualitative sense that is. A big bang is still possible, but (a) how 'big' that 'bang' was we cannot know, (b) it is possible that there were many 'bangs' (in space and/or time), and (c) there is no need to invoke 'bangs of creation' or (d) to consider 'bangs' that originate out of a 'point' in space. Shells can collide (subsections \ref{sect_phasefunction} and \ref{sect_collapse}), or put differently, a universe not necessarily evolves in a synchronous way. This possibility adds an important twist to cosmological models, since RW models (these belong to a class which we call synchronous models) do not allow for collisions. Colliding shells provide a natural mechanism to make filamentary structures\footnote{In this paper we remain in the spherical symmetry. The structures are thus spherical shells.}. Collisions also mark singularities in the metric. These do not invalidate the geometrical theory of gravitation, of course, but their occurrence shows the danger in assuming that one metric can be employed, without limits in space and time, to describe the universe. We argue that there is an important difference between the validity period of a (geometrical) cosmological model and the real age of the universe. We must acknowledge that the latter may be impossible to determine, in contrast with the age of the structures in it, which can be dated. If we consider the consequences of collisions together with the ideas of the previous paragraph, one can envisage a universe as comprised of many smaller entities, somewhat like a raspberry or a pomegranate. We consider a rather general class of specific cosmological models, but must admit, of course, that we cannot be exhaustive (subsection~\ref{sect_asynchronous}). Also models that have infinite extent or that are not the interiors of black holes are part of this class, but are not the subject of this paper. Barring these models, our models have an outermost shell, the boundary, beyond which there is no matter. The RW models are 'almost everywhere' identical to the synchronous models (subsection~\ref{sect_synchronous}), the exception being that they differ from these on the boundary, which the RW models simply don't have. The Hubble parameter is dependent on the viewing direction, and we define $H_{\rm o}$ to be the current value as measured in the direction perpendicular to the center of the universe. This definition is consistent with the original definition in the standard cosmological models. We consider only models that realise $H_{\rm o}$ somewhere and sometime (subsection~\ref{sect_H0}), and classify them in some loose sense accordingly (subsection~\ref{sect_asynchronous}). In subsection~\ref{sect_behaviourX} we also consider another, rather technical, classification that has to do with the behaviour of the radial metric coefficient at the boundary. It is important in the calculation of orbits. Contrary to the standard models, we have to cope with plane orbits instead of radial orbits, which is the subject of section~\ref{sect_orbits}. We devote a special section (section~\ref{sect_light}) to light, since it can be argued that photons are the only particles that are essentially not comoving and thus can be considered to be traveling on a cosmological scale. In subsection~\ref{sect_orbitbound} we examine the behaviour of the orbits in the vicinity of the boundary. We show that no particle can cross the boundary. Upon arrival they are reflected back into the universe. The boundary therefore effectively acts as a gravitational mirror. Unlike an ordinary mirror, however, one effectively flies into the mirror image upon arrival at the boundary, causing a spinor-like effect. This mirror-boundary could have far-reaching consequences. To pick one, it could be that our universe looks larger than it is. Indeed, we may find ourselves as if we were inside a spherical mirror. For that mirror to have effective observational consequences one needs (light travel) time, however (and transparent space of course). We discuss this and other rather intriguing aspects of the boundary in subsection~\ref{sect_physnatrb}. If our universe is confined inside a Schwarzschild horizon, we must also consider 'the mother universe'. In particular, we must consider material that falls into our universe from there. The connection between the Schwarzschild-$\Lambda$ metric of our mother universe and the metric of our universe is elucidated in appendix~\ref{sect_Novikov}. We employ Novikov coordinates, which are the coordinates of an observer who moves in a swarm of radially infalling shells, and who determines his radial position by the shell that passes by. These coordinates give rise to a diagonal metric, but are nevertheless largely unused because of the metric coefficients which are fairly complicated in their original form. The function 'cyc' however hides the mathematical details, and the resulting metric turns out, in hindsight not surprisingly, to be of the form of the metric of our universe, but it is one that is, in contrast, collapsing. Additional great features of the Novikov coordinates are that they can be easily interpreted and there is no horizon singularity. Hence it is possible to treat the infalling particle problem, since spacetime is not treated differently inside or outside the Schwarzschild radius. This is outlined in section~\ref{sect_beyond}, with the aid of the formalism developed in subsection~\ref{sect_NoviSchwarz}. In section~\ref{sect_beyond} we also discriminate between the Schwarzschild sphere of the universe and the boundary of the universe. The picture emerges, at least in the ideal spherical and ordered case without colliding shells, of a kind of 'breathing' universe in the spacetime that does not belong to the universe and that is 'interior' to the Schwarzschild radius.\footnote{Strictly spoken, there should be nothing left between the edge of the universe and its Schwarzschild radius, but we keep the possibility open that, with the help of non-spherical evolutions in the past, there are some leftovers from a previous expansion phase. In any case, whatever there is, it continuously falls back into the universe.} The maximum expansion of the universe occurs at the very moment when also the universe fills its Schwarzschild sphere. At all other times, it does not, and there is (empty) space-time between the boundary of the universe and the Schwarzschild horizon, space-time that can be described with a Novikov metric. In this picture, it therefore is not correct to ascertain that an expanding universe 'creates' space along with its expansion. Our universes simply expand into the preexisting space-time interior to their Schwarzschild spheres. The above picture also points to an intriguing possibility that can occur when a black hole with a universe inside is in a state in which its universe reaches the Schwarzschild radius. In section~\ref{sect_limit} we show that, if at that very moment of maximum expansion the shells at the boundary also collide, it is possible that the boundary, and hence the Schwarzschild shell, is crossed from the inside to the outside. Unlikely as this may be globally, it could be realized locally (thereby breaking the spherical symmetry of course). This could be relevant for 'ordinary' black holes, which are most likely fast rotating and described by a Kerr metric. This metric however is locally, at the poles, the same as the Schwarzschild metric, and therefore emission at the poles could be a possibility, thereby enabling bipolar outflows. In section~\ref{sect_anom_red} we link this mechanism also to H. Arp's anomalous redshifts, since material that manages to escape from a black hole must climb out of a rather formidable gravitational well. Barring the cases of chance alignments, the anomalous redshifts are therefore not cosmological but gravitational according to this mechanism.\footnote{We do not advocate that all Arp's cases are of this nature. Surely there can be chance alignments. Hence statistics on Arp's catalog is not very relevant to the point made in this paper.} In standard cosmologies, any point can be the center of the universe, any light ray can therefore be a radial one, and flux calculations are straightforward by considering the surface of a sphere centered at the origin. Not so in this case, and in appendices \ref{sect_angdistance} and \ref {sect_surface brightness_flux} we set up the necessary machinery to calculate the magnitude as a function of redshift (section~\ref{sect_light}), which is needed to fit models with observations. We find, of course, that the relation depends on the direction of the line-of-sight. Though a fit to observations is way beyond the scope of this paper, we show in section~\ref{sect_comp_obs} that the presence of a phase function suffices to fit current observations without having to invoke dark energy. We also compare some of our models with the consensus model, and find that we can reproduce it, of course, but also with $\Lambda=0$.\footnote{Since this paper includes the case $\Lambda\ne0$ and RW models are part of the class we consider (almost everywhere), a satisfying fit to observations is a trivial matter anyway.} In section~\ref{sect_darkenergy} we briefly comment on dark energy. Since dark energy and $\Lambda>0$ are essentially the same thing observationally, in our models a $\Lambda\ne0$ is a property of the mother universe, of which we never will know anything except via whatever falls from there into our universe. In section~\ref{sect_primBH} we speculate about the formation of galaxies. Not much more can be done, since our models have a history before the latest expansion phase that we are in now. We argue that in our paradigm one should seriously consider that at least some black holes are primordial to our expansion phase, either by infall from our mother universe, or by growth in a collapse phase. Indeed, we show that black holes can easily gobble up a lot of material if the collapse around it is strong enough. The well-established relation between central black holes and the mass of the halo for spiral galaxies suggests that at least spirals are basically giant accretion disks around massive black holes, formed in a (local) collapse phase. The picture of our universe as a black hole in a mother universe, together with our universe comprising numerous super massive black holes that are the seeds of new universes, is logically consistent and therefore rather appealing. As to the dark matter, in section~\ref{sect_darkmatter} we pick up the old idea that it is made up of neutrino's, but stress that the analysis is valid for any particle that obeys a Fermi-Dirac or Boltzmann distribution. For this to work, the dark matter must be massive enough, cold enough, and dense enough, all of which are problematic properties in standard cosmology because they would close the universe. Not here, on the contrary: our models are closed by design. We show in section~\ref{sect_neut_BH} that the remnant of a gravitational collapse that forms or grows a primordial black hole is a nearly constant density sphere of trapped dark matter particles. For a Milky Way type black hole, that sphere easily may contain as much mass as the black hole itself within about 5 pc, which is not contrary to observations of our Milky Way black hole environment. For larger black holes, the dark matter envelope can be very massive depending on its extent, and therefore the inferred black hole masses can include a very important contribution from the dark matter envelope. On a larger scale, the flat rotation curves of spirals show that the dark matter distribution must have been isothermal at the moment in the expansion when they were 'locked in' into the galaxies (section~\ref{sect_intro_neutrinohalo}). Just like the isotropy of the CMB, this points to the scenario that our universe must be much older than the latest expansion phase, such that it was able to also produce a isotropic and isothermal cosmic dark matter background. It should be emphasized, though, that the structures we now see need not nearly be as old. In section~\ref{sect_simple_model} we produce a simple spherical model which includes the essential ingredients of a Milky Way type galaxy with a flat rotation curve. In the final section, we revisit Mach's principle in the light of our universes. We show that inertial mass can be explained as a kind of stiffness that every massive particle acquires as a consequence of the attraction by the cosmic dark matter background. | 18 | 8 | 1808.10278 |
1808 | 1808.02830_arXiv.txt | We simulate star formation in two molecular clouds extracted from a larger disc-galaxy simulation with a spatial resolution of $\sim 0.1$\,pc, one exiting a spiral arm dominated by compression, and another in an inter-arm region more strongly affected by galactic shear. Treating the stars as `sink particles', we track their birth angular momentum, and the later evolution of their angular momentum due to gas accretion. We find that in both clouds, the sinks have spin vectors that are aligned with one another, and with the global angular momentum vector of the star cluster. This alignment is present at birth, but enhanced by later gas accretion. In the compressive cloud, the sink-spins remain aligned with the gas for at least a free fall time. By contrast, in the shear cloud, the increased turbulent mixing causes the sinks to rapidly misalign with their birth cloud on approximately a gas free-fall time. In spite of this, both clouds show a strong alignment of sink-spins at the end of our simulations, independently of environment. | The densest parts of molecular clouds -- molecular cloud cores -- collapse under their own gravity, protected from any external radiation by an envelope of dust. These are the sites where new stars are born \citep{andre2013}. However, the fine details of this process of star formation remain unclear. One potentially powerful constraint on formation channels comes from the alignment or non-alignment of star spins inside star clusters. The latest theoretical models find that the angular momentum of molecular cloud cores depends primarily on small scale turbulence, though breaking due to magnetic fields that can also reduce their angular momenta \citep{2007ARA&A..45..565M}. The spin-alignment of stars in stars clusters depends, therefore, on what fraction of the kinetic energy of the host cloud is in rotational versus turbulent pressure support \citep[e.g.][]{Corsaro2017}. The observational picture of spin alignment, however, remains rather murky. Early work from \cite{Jackson&Jeffries2010} used estimates of the radii of stars, combined with spectroscopic rotation measurements, to estimate spin-alignment in the Pleiades and Alpha Per star clusters (with masses, $\sim800$\,M$_\odot$ and $\sim 350$\,M$_\odot$, respectively). They found no evidence for strong alignment, a result that persists even given new and larger estimates of the radii of M dwarf stars in Pleiades \citep{2018MNRAS.476.3245J}. By contrast, \cite{2018A&A...612L...2K} recently used a similar technique to rule out an isotropic distribution of spins in the Praesepe cluster (with stellar mass 500-600\,M$_\odot$), but see \citealt{2018MNRAS.tmp.1301K} for a discussion of some of the difficulties inherent in this measurement. Finally, using a different technique that measures stellar spins using astroseismology, \citet{Corsaro2017} found that stars in the open clusters NGC 6791 and NGC 6819 (with stellar masses $\sim$5000 M$_{\odot}$ and $\sim$2600 M$_{\odot}$, respectively) appear to have highly aligned spin vectors. Using simulations with a spatial resolution of 0.002\,pc \citep[see][]{Lee&Hennebelle2016}, they show that such high alignment may be achieved if $>50$\% of the kinetic energy of the natal cloud is in rotation. (Note that all of these studies can be compatible if higher mass star clusters show stronger spin-alignment than lower mass clusters, or if astroseismology is a more reliable probe of spin-alignment than spectro-photometric techniques.) As discussed above, from a theoretical perspective the spin-alignment of stars in star clusters depends sensitively on the turbulent velocity field of their birth clouds. This, in turn, depends on the galactic environment of the cloud \citep{Renaud2013,rrr2015}. Gas clouds formed in spiral arms are shocked and compressed, leading to the formation of more stars, while if a cloud forms in an inter-arm region, it is more dominated by shear, with less star formation. In this Letter, we study whether molecular clouds extracted from a spiral (compressive) or inter-spiral (shear) region lead to systematic differences in the spin-alignment of their stars. We run our simulations at a resolution of 0.1 pc, which is around 2 orders of magnitude smaller than the half light radius of 2 typical open clusters NGC 6791 and NGC 6819 (10\,pc and 7\,pc, respectively). We use sink particles to model the formation of stars. Our goal is to determine whether spin-alignments are ubiquitous, given realistic initial conditions for the birth-cloud, or whether they depend on environment. | In this Letter, we have studied the effect of the galactic environment in the transference of angular momentum between a parent gas cloud and its molecular cores. Our results suggest that, at creation, the spin of a sink particle follows the angular momentum of the gas in its local galactic environment. Even in two very different galactic environments, we find that the sinks are aligned both with the global angular momentum of the cluster, and with the average angular momentum of the stars. Our results require confirmation from simulations of a larger number of molecular clouds at higher resolution and including the physics of magnetic fields and stellar feedback. However, the fact that we find that star spins are strongly aligned in two very different galactic environments suggests that star-spin alignments may be ubiquitous. This has interesting implications for the formation of massive stellar binaries and their gravitational wave emission. Our results suggest that such binaries are much more likely to be spin-aligned, leading to an incorrect inference of their stellar remnant properties if non-aligned templates are used \citep[e.g.][]{2017JKPS...70..735C,2018JCAP...03..007C}. | 18 | 8 | 1808.02830 |
1808 | 1808.00835_arXiv.txt | The concept of cosmological inflation~\cite{Starobinsky:1980te, Sato:1980yn, Guth:1980zm, Linde:1981mu, Albrecht:1982wi, Linde:1983gd} has become a vital part of the standard cosmological model, in particular, thanks to providing a mechanism for generating of primordial density perturbations~\cite{Starobinsky:1979ty, Mukhanov:1981xt, Hawking:1982cz, Starobinsky:1982ee, Guth:1982ec, Bardeen:1983qw}. However, after almost four decades of theoretical pursuit inflation remains a very general theory with no direct link to the Standard Model of particle physics (with a notable exception of so-called Higgs inflation \cite{HI1,HI2,HI3,HI4}). In particular, there is no universally accepted mechanism of reheating, {\it i.e.} the transition between inflationary era and radiation domination era. Various plausible scenarios for reheating \cite{BT,KLS}, well embedded in the framework of quantum field theory have been thoroughly investigated (see, {\it e.g.}, \cite{Bassett:2005xm,Allahverdi:2010xz} for a review). One appealing possibility relies on mode amplification of quantum fluctuation and particle production that can occur when the homogeneous part of the inflaton field oscillates around the minimum of its potential \cite{PRH2,PRH3}. Quite recently, it has been understood that in some models such oscillations can excite fluctuations of the inflaton field to the extent that the Universe starts expanding as radiation dominated. This mechanism, called self-resonance \cite{MA,AEF,AEFFH}, offers an economical and elegant exit from inflation in some inflationary models, greatly reducing theoretical uncertainty associated with reheating \cite{RE}, For a given model, one can numerically investigate self-resonance with lattice simulations, but this is both memory- and time-consuming and, so far, has been done only for a few inflationary models, see, {\it e.g.}, \cite{AL0,AL}. However, before undertaking full nonlinear lattice simulations one can approach the problem at linear order in perturbations and, using Floquet theory, predict whether self-resonance can lead to efficient reheating (see, {\it e.g.}, \cite{FE} for a review). In this letter, we utilize {\it Encyclopaedia Inflationaris}, a comprehensive review of single-field models of inflation \cite{EI}, to identify models which are both consistent with observational data (including the 2018 Planck data release \cite{PC}) and admit efficient self-resonance as the mechanism for reheating. To this end, we employ Floquet analysis of inflaton perturbations. The letter is organized as follows. In Section \ref{sec:two} we briefly present the self-resonance mechanism and briefly review Floquet theory. In Section \ref{sec:four}, we present the overview of single field inflationary models together with their prospective Floquet analysis. We draw our conclusions in Section \ref{sec:five}. | \label{sec:five} We performed Floquet analysis of the evolution of the perturbations of the inflaton after the end of inflation in a number of theoretically justified and observationally favored single-field inflationary models considered by other authors. We showed that these perturbations can be unstable and lead to to the fragmentation of inflaton condensate in six of investigated models. However, we found that, in addition to the already known example of $\alpha$TI with $n\neq1$, only in case of KKLT inflation with $p\neq 2$, self-resonance can be a good scenario for reheating, {\it i.e.} describe the transition of the equation of state of the Universe close to that of radiation. In the remaining cases, the instability leads to creation of long-lived oscillons, for which the typical equation of state is that of non-relativistic matter. Therefore, in those cases, other scenario of reheating are necessary. Taken together, our findings suggest that efficient reheating {\it via} self-resonance, although possible in principle, is rather rare among inflationary models. \subsubsection*{Acknowledgements} The authors are supported by grant No.\ 2014/14/E/ST9/00152 from the National Science Centre (Poland). \begin{figure} \centering \begin{tabular}{cc} \includegraphics[width=0.48\textwidth, height=0.25\textwidth]{HOverOmegaPlot.pdf} & \includegraphics[width=.48\textwidth, height=0.25\textwidth]{dtHOverHOmegaPlot.pdf} \end{tabular} \caption{\it Values of $H/\omega$ and $\dot{H}/H\omega$ for the first 15 oscillations of the inflaton field around the minimum of the potential for selected models considered in Sections \ref{sec:small} and~\ref{sec:large}. \label{fig:ex1}} \end{figure} | 18 | 8 | 1808.00835 |
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1808 | 1808.03836_arXiv.txt | The discovery of GW170817 with gravitational waves (GWs) and electromagnetic (EM) radiation is prompting new questions in strong-gravity astrophysics. Importantly, it remains unknown whether the progenitor of the merger comprised two neutron stars (NSs), or a NS and a black hole (BH). Using new numerical-relativity simulations and incorporating modeling uncertainties we produce novel GW and EM observables for NS-BH mergers with similar masses. A joint analysis of GW and EM measurements reveals that if GW170817 is a NS-BH merger, $\lesssim 40\%$ of the GW parameters are compatible with EM observations. | 18 | 8 | 1808.03836 |
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1808 | 1808.08912_arXiv.txt | The direction of the axis of the orbital motion of the merging binary neutron stars in the GW170817 event coincided with that of the apparent superluminal jet, which produced the short hard gamma ray burst (SHB) 170817A. It supports the local value of the Hubble constant provided by standard candle Type Ia supernovae, $H_0\!=\!73.24\! \pm\! 1.74\, {\rm km/s\, Mpc}$, which differs by 3$\sigma$ from the cosmic value $H_0\!=\!67.74\!\pm\! 0.46\,{\rm km/s\ Mpc}$ obtained from the cosmic microwave background radiation by Planck assuming the standard $\Lambda$CDM cosmology. The measured superluminal motion of the jet also allows critical tests of the assumed production mechanism of SHBs in general and of SHB170817A in particular. | 18 | 8 | 1808.08912 |
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1808 | 1808.03694_arXiv.txt | We study the formation of classical singularities in Generalized Brans-Dicke theories that are natural extensions to Brans-Dicke where the kinetic term is modified by a new coupling function $\omega(\varphi)$. We discuss the asymptotic limit $\omega(\varphi)\rightarrow\infty$ and show that the system generically does not approach General Relativity. Given the arbitrariness of $\omega(\varphi)$, one can search for coupling functions chosen specifically to avoid classical singularities. However, we prove that this is not the case. Homogeneous and spherically symmetric collapsing objects form singularities for arbitrary coupling functions. On the other hand, expanding cosmological scenarios are completely free of Big Rip type singularities. In an expanding universe, the scalar field behaves at most as stiff matter, which makes these cosmological solutions asymptotically approach General Relativity. | In a recent letter, Alexander et al.~\cite{alexander2016} investigate possible Higgs-like mechanisms to account for the emergence of gravity in extensions to General Relativity (GR). The main idea is that symmetry restoration (breaking) could ``turn off'' (``turn on'') gravity. The authors argue that a possible realization of this scenario is to identify the scalar field in Brans-Dicke (BD) theory with the Higgs-like field that undergoes symmetry restoration. At low curvature configurations the system must conform to GR but at high curvature regime a phase transition can drastically modify the dynamics. In particular, one might wonder if this type of Higgs-like mechanisms could dynamically avoid classical singularities. BD theory is an attempt to incorporate Mach's Principle in a relativistic gravitational theory. This pioneer work by C. Brans and R. H. Dicke \cite{brans1961} can be seen as a scalar-tensor gravity theory where the extra scalar field is non-minimally coupled to the Ricci scalar. This coupling introduces an effective gravitational strength inversely proportional to the scalar field, $G_{eff} \sim \varphi^{-1}$. Recently, BD received much attention due to its phenomenological applications in cosmology~\cite{hrycyna2014,hrycyna2013,papagiannopoulos2016,alonso2016,roy2017}. The kinetic term associated with $\varphi$ carries a constant dimensionless parameter $\omega_{BD}$. In order for BD to be consistent with solar system astronomical experiments, the BD parameter must be large ($\omega_{BD} > 500$ see~\cite{will1993}). Effectively, this constrains BD to be observationally indistinguishable from GR in the weak field limit. There are black hole solutions in BD~\cite{bronnikov1973,bronnikov1999,scheel1995,campanelli1993}, even though not with exactly the same features as in GR, and an analogue version of Birkhoff's theorem that proves that a static, spherically symmetric, asymptotically flat, vacuum solution in the Jordan frame is uniquely characterized by the Schwarzschild solution~\cite{faraoni2018}. Thus, it is certain that BD has singular solutions. A straightforward generalization of BD is to promote $\omega$ to be a function of the scalar field as it is the case of generalized Brans-Dicke theories (GBD). In fact, BD can be seen as a particular example of GBD where the coupling function is fixed to be a constant. GBD was first introduced by Nordvedt \cite{nordtvedt1970}, where he analyze the post-Newtonian corrections to weak-field regime. Almost simultaneously, Wagoner~\cite{wagoner1970} also studied the weak-field limit considering the predictions for classical tests as light propagation and perihelion shift effects and the observable differences between the scalar and tensor components in gravitational waves. Our goal in the present work is to analyze if GBD can dynamically avoid the formation of singularities. For that we shall discuss collapsing and cosmological solutions. In the next section we briefly describe GBD theories and the conditions for approaching the GR regime. In section \ref{MCGS} we discuss the necessary matching conditions and the geometrical setting for our dynamic system. In section \ref{CCO} and \ref{COSMO} we analyze, respectively, the collapsing and cosmological scenarios and in \ref{Conclusions} we conclude with some final remarks. | Generalized Brans-Dicke theories are natural extensions to Brans-Dicke original proposal that maintains the same non-minimal coupling between the curvature and the scalar field while introducing a new coupling function to its kinetic term. The main motivation is to allow the Brans-Dicke parameter to vary in different gravity scenarios, such as, between primordial universe and solar system dynamics. In the present work we have studied the formation of classical singularities in GBD. Given the arbitrariness of the coupling function, one could argue that with an adequate choice of $\omega(\varphi)$, in principle, it would be possible to dynamically avoid classical singularities. The simplest scenarios are homogeneous and isotropic spatial section described by the family of FLRW metrics. Our analysis depends only on two physically motivated hypotheses. Solar system experiments and cosmological observations show that GR is in good agreement with experimental data\cite{will1993,ade2015.1,ade2015.2}. Thus, we assume initial conditions that mimic a GR regime, namely, small scalar field velocities $\left(|H_{\varphi}|\ll 1\right)$ and large coupling function ($\omega\gg1$). We impose boundary conditions for the coupling function such that the effective gravitation strength, $G_{eff} \sim \varphi$, is non-negative and satisfies the conditions given by figure~\eqref{figcouplingfunc}. We have shown that spherically symmetric and homogeneous collapsing objects generically form singularities. Case I-b can avoid the singularity but the solution is fine-tuned and not physically realizable. Indeed, it seems reasonable that less symmetric solutions might relax this fine-tuned condition for singularity avoidance; we shall analyze carefully this possibility in a future work. Contracting FLRW cosmological solutions are formally identical to the collapsing cases. Thus, the cosmological singularities that plague GR are also present in GBD. On the other hand, expanding cosmological scenarios are completely free of Big Rip-type singularities. We have shown that in an expanding universe the scalar field behaves at most as a stiff matter type fluid, which makes the system asymptotically approach GR. | 18 | 8 | 1808.03694 |
1808 | 1808.03377_arXiv.txt | In 2015 the Institute of Physics and Astronomy of the Universidad de Valparaíso in Chile received as a donation the Bochum 0.61-meter telescope. Here we preset the ongoing project to convert this senior member of La Silla Observatory to modern standards aiming at performing state-of-art science, as well as teaching and outreach. Firstly, the site characterization was performed in order to verify the observing conditions. The preliminary results were auspicious in relation to the nights available for observation. In early 2016 began the transfer work form La Silla Observatory to the new site of operations. The actual status of the telescope was analyzed and an upgrade plan was proposed to make it usable remotely using a web-based telescope control system developed in Chile by ObsTech SpA. Future upgrade and scientific collaboration will be discussed based on the site characterization and technical studies regarding the potential for new instrumentation. | \label{sec:intro} % The Bochum 0.61-meter telescope, installed on La Silla in 1968, was the first national telescope in ESO's history through a trilateral agreement between ESO, the Deutsche Forschungsgemeinschaft (German Research Foundation) and the University of Bochum. In the early 70s this telescope was one of the so-called ``beacon telescopes'' in the mountain, given its remarkable photo-spectrograph and also is remembered to be one of the first telescope to detect unexpected evolution phenomena from supernovae SN 1987A, still referred to as ``Bochum event'' in scientific literature \cite{hanuschik_1988}. The new professional telescope now of the Universidad de Valparaíso (including its original mounting and copula) was driven by a dream of Dr. Nikolaus Vogt, Professor and Researcher of the Institute of Physics and Astronomy (IFA) of the Universidad de Valparaíso (UV) in Chile, who made the arrangements with the University of Bochum to make the donation of this astronomical professional instrument, that was not being used anymore on La Silla, to the UV. Studies were carried out to find the observationally most optimal cite in the Valparaíso region. Finally, the ideal location was found to be in Pocuro at Calle Larga town (around 2 h from Valparaíso city). At Calle Larga already exist active amateur astronomer community called \textit{Asociación Astronómica de Aconcagua} who founded the Observatorio Pocuro (including a set of small scale telescopes). After a collaboration agreement between the UV and Calle Larga authorities, the Bochum 0.6-meter telescope was placed in its new home at Observatorio Pocuro (see Fig. \ref{fig:pocuro_telescope}). In early 2016 the transfer preparation work began to transport the Bochum 0.6-meter telescope from La Silla Observatory to the new site of operations in the Observatorio Pocuro at Calle Larga. The transfer process involved management and logistic around safely transportation of the cupola and telescope. A new building was constructed at the Observatorio Pocuro in order to set up the cupola and the telescope properly. Our goal is to modernize the Bochum telescope by upgrading the Telescope Control System (TCS) in order to replace the DOS-based system currently installed, along with the old power control electronics and encoders. The actual status of the telescope was analyzed and an upgrade plan was proposed to make it usable remotely using a web-based telescope control system developed in Chile by ObsTech SpA \cite{1mSilla}. The aim of the renovated Bochum telescope is multipurpose. Our project explores the huge potential of the small previously abandoned telescope for teaching, science and outreach purposes. This potential will be discussed and future plan will be presented. \begin{table}[h!] \begin{center} \label{tab:table1} \begin{tabular}{|l|c|} % \hline \textbf{Enclosure} & Classical dome \\ \hline \textbf{Type} & Photometric telescope \\ \hline \textbf{Optical design} & Cassegrain reflector \\ \hline \textbf{Diameter Primary M1} & 0.61 m \\ \hline \textbf{Material Primary M1} & Low expansion Silica \\ \hline \textbf{Diameter Secondary M2} & 0.15 m \\ \hline \textbf{Material Secondary M2} & Low expansion Silica \\ \hline \textbf{Mount} & Equatorial cross-axis mount \\ \hline \textbf{First Light date} & 7 September 1968 \\ \hline \end{tabular} \caption{Bochum 0.61-meter telescope summary.} \end{center} \end{table} \begin{figure}[H] \centering \includegraphics[width=.6\textwidth]{telescopio.png} \caption{The Bochum 0.61-meter telescope at Observatorio Pocuro.} \label{fig:pocuro_telescope} \end{figure} | The revival process of the Bochum 0.61 meter telescope, a former abandoned member of La Silla Observatory, was documented. This was a project with great challenges since its very inception which has nonetheless given us a valuable learning experience. Currently we are working on the renewal of the Bochum telescope in order to give give a ``Second life'' to this abandoned telescope and offer the students as well as the researchers of Universidad de Valparaíso a solid observational potential, as well as teaching and outreach opportunities. One of the applications of the modernized Bochum telescope will be teaching both undergraduate and graduate students of the UV in the field of observational astronomy with the possibility to make observations on the site as well as remotely, and to familiarize with the procedures regarding telescope operations. Complementary to teaching our 0.6-m telescope will be used for research projects of the members of our Institute and will collaborate on several survey campaigns. Finally, but not less important, our Bochum telescope will have the strong outreach mission, we will offer 20\% of the time to the local community and the Observatorio Pocuro which will use it to bring the local community closer to the wonders of astronomical observation and science in general. | 18 | 8 | 1808.03377 |
1808 | 1808.04885_arXiv.txt | This is an opinion-piece based on a talk given at the Summer 2017 Serendipities in the Solar System (``Ip-Fest'') meeting in Taiwan\footnote[2]{A celebration of Wing Ip's 70th birthday, held at the Institute of Astronomy, National Central University, Taiwan, July 2017; published in 2018 as Astronomical Society of the Pacific Conference Series Volume No.~513, page 33}. I begin with an overview of the new solar system, then discuss modelling attempts, then the distribution of optical colors as a proxy for the distribution of compositions, and I end with remarks about Wing Ip. | Three cometary reservoirs exist (see Figure \ref{schematic} and Jewitt et al.~2009 for a review): 1) The smallest is the asteroid-belt itself which, despite its high radiation equilibrium temperatures, preserves bulk ice (see also Snodgrass et al.~2017). In the main-belt comets, ice is shielded from the Sun by refractory mantles and only occasionally sublimates when exposed by impact or other surface disturbances. The duty cycle (ratio of ``on'' time to total elapsed time) is $< 10^{-4}$. There are about 10$^6$ asteroids larger than a kilometer in scale. The fraction containing ice is uncertain, but may approach unity in the outer-belt. 2) The largest reservoir is the Oort cloud, a spherical swarm of vast extent ($\sim10^4$ to $10^5$ AU scale) from which the long-period comets are supplied (see also Rickman 2014). The number of Oort cloud comets is uncertain, but is probably in the 10$^{11}$ to 10$^{12}$ range. Their combined mass is likely in the 1$M_{\oplus}$ to 10$M_{\oplus}$ range (1$M_{\oplus}$ = 6$\times$10$^{24}$ kg), but could be a little smaller or much larger, depending mainly on the unmeasured properties of the inner Oort cloud. 3) The third cometary reservoir is the Kuiper belt (a.k.a.~the trans-Neptunian belt), source of most short-period comets. The Kuiper belt holds at least 10$^9$ objects larger than 1 km and 10$^5$ objects larger than 100 km, with a combined mass not more than $\sim$0.1$M_{\oplus}$. The three reservoirs sample products of low temperature accretion in different radial regions of the protoplanetary disk. The main-belt comets likely formed in-place, close to the snow-line in the epoch of accretion, although other formation locations have been suggested. The Oort cloud comets probably formed in the middle solar system, where we now find the giant planets, and were subsequently hurled outwards by near-miss gravitational scattering from the planets as they grew. Most were lost to the interstellar medium, never to be seen again, but maybe 1\% to 10\%~of the comets were deflected on the way out by the combined gravitational effects of passing stars and the galactic tide. They became trapped in long-lived, weakly-bound heliocentric orbits with perihelia far outside the planetary region (e.g.~Dones et al.~2015). The Kuiper belt contains a mixture of objects that we think were scattered into place from interior source regions (the so-called dynamically hot populations) with objects that likely formed in-place (the cold classical objects). The reservoirs thus offer a kind of radial stratigraphy of the solar system, by providing icy products accreted at temperatures $\sim$140 K (main-belt comets) to $\sim$40 K (cold classicals). \articlefigure{Jewitt_F1.pdf}{schematic}{Schematic diagram showing the relationships believed to exist between various small-body solar system populations. Acronymns in the Figure indicate different sub-populations: LPCs = Long-period comets, HTCs = Halley-type comets, JFCs = Jupiter family comets, dJFCs = dead (or dormant) Jupiter family comets. Damocloids are probably dead (or dormant) HTCs. Numbers on the right indicate that, while the reservoirs survive for the age of the solar system, bodies outside the reservoirs meet a rapid demise. Bodies scattered by the giant planets have median lifetimes $\sim$10$^7$ yr while those in the cramped quarters of the terrestrial planet region have dynamical lifetimes $\sim$0.5 Myr (and physical lifetimes considerably shorter). All comets displaced from their reservoirs meet similar fates.} Study of the Oort cloud is limited by its vast size, which renders its constituent comets too faint to be detected from Earth. We are forced to infer the properties (and existence!) of the cloud from comets that have left it (the problem is analogous to trying to figure out the size and layout of the parking structures at LAX by only counting the number of cars leaving the airport). Study of the main-belt comets is limited by the small number (only a handful) of known examples (the analogy is trying to figure out the parameters of LAX parking by staring at just one or two parked cars), although there is hope that intensive future observations will change this circumstance. In contrast, the Kuiper belt is eminently observationally accessible, and the physical and dynamical properties of large numbers of Kuiper belt objects can be measured directly. Indeed, about 2000 Kuiper belt objects are known at the time of writing and from them the incredible and unexpectedly complicated architecture of the outer solar system has been revealed. \textit{The special feature of the Kuiper belt, and the reason that its discovery has revolutionized planetary science, is that it is near-enough for its contents to be studied directly}. Measurements of the orbits of Kuiper belt objects have provided the food for modellers interested in the dynamical aspects of the origin and evolution of the solar system. The key observational finding was that the abundance of resonant objects (particularly the 3:2 ``Plutino'' population) is too large to be a result of chance. Instead, the resonant populations are best understood as a consequence of the radial migration of Neptune during the formation epoch (Fernandez and Ip 1984, Malhotra 1995). If Neptune's orbit changed size, then so must have the orbits of the other planets, with potentially profound dynamical consequences that depend on the degree and the rates (and smoothness) of the migration. Extreme effects might result if the giant planets were driven to cross mean-motion resonances with each other, as described in the famous Nice model and its variants, derivatives and follow-ons. Significantly, such effects might include the capture of objects scattered from the Kuiper belt into dynamical niches including the Jovian Trojans (Morbidelli et al.~2005), the irregular satellites (e.g.~Nesvorny et al.~2014) and even the asteroid belt (Levison et al.~2009). Most recent solar system models are numerical N-body simulations, relying on fast computers to represent complex dynamical systems. The biggest strength of the numerical models is their flexibility; their many parameters can be readily adjusted to fit many (but not all) of the known properties of the solar system\footnote{An obvious exception is the cold-classical Kuiper belt which, with its low mass, narrow inclination distribution, high abundance of binaries and sharp outer edge, appears in the Nice model and its variants as an ad-hoc addition, not as a natural consequence of planetary migration into a pre-existing planetesimal belt (c.f.~Fraser et al.~2014).}. Their biggest weakness is exactly the same thing - their flexibility, which is so great that they struggle to offer any firm, observationally testable predictions. The result is scientifically frustrating in the sense that interesting models come and go (the Nice model, the modified Nice model, the jumping Jupiter model, the Grand Tack model, models in which extra giant planets are added to the solar system then allowed to escape, models in which extra components are added to the asteroid belt and then dynamically destroyed) but nobody can tell whether they describe the real world, or just a model-world. Consider the late-heavy bombardment (LHB) of the Moon at 3.8 Gyr as a case in-point. The Nice model was originally presented as a solution to the puzzle of why the LHB was delayed for 800 Myr following the Moon's formation (Gomes et al.~2005). It did this rather ingeniously, by slowly driving the planets towards a resonant instability, using torques exerted by a Kuiper belt selected to have the ``right'' mass and the ``right'' separation from the outermost planet. However, from the start, the existence of the LHB was doubted by the community best equipped to judge it (geologists and geochemists who studied the Apollo lunar samples, see Hartmann 1975), a fact that was forgotten, ignored or disputed by the dynamicists. Improved measurements of lunar rocks (especially the highly refractory and isotopically revealing nuggets called ``zircons'') have strengthened the alternate explanation, namely that the impact flux declined smoothly over a long period of time, not in a spike-like LHB (see Zellner 2017, for a recent review). In other words, the very problem that the Nice model was proposed to solve has evaporated. The response of the Nice modelers has been to simply change a few assumptions about the initial conditions of the solar system in order to push the LHB closer to the formation epoch (Morbidelli et al.~2017). Having no evidence for the LHB could simply mean, in the context of the model, that resonance crossing occurred so close in time to the formation of the planets that it is indistinguishable from the heavy cratering flux in the accretion phase. This might be true, of course, but it might also be true that the LHB never occurred and that the Nice model, as constructed, describes something that didn't happen. The fundamental problem is that the numerical models can be tuned to provide a wide range of outcomes, but they don't offer the means to test the veracity of those outcomes. | \begin{itemize} \item Discoveries in the Kuiper belt have prompted a stream of increasingly elaborate, but largely untestable, numerical models of solar system evolution. While it is clear that we have a much improved appreciation for the complexity of the solar system, there is also an undeniable feeling that the models have passed the point of diminishing scientific returns. Something must change if we are to make real progress. \item Kuiper belt populations, both dynamically hot and cold, include large fractions of objects with ultrared (B-R $>$ 1.6, $S' >$ 25\%/1000\AA) colors, tentatively interpreted as a marker for cosmic ray-processed complex organics. About 2/3rds of cold-classical KBOs and 1/3rd of the hot populations are ultrared. \item High perihelion ($q \gtrsim$ 8 to 10 AU) Centaurs closely resemble the Kuiper belt hot population by including a $\sim$1/3rd ultrared matter fraction in their optical color distribution. \item Smaller perihelion Centaurs are depleted in ultrared matter, probably as a result of a temperature-dependent evolutionary effect (most likely ballistic resurfacing, given the coincidence between the critical distance for the onset of Centaur activity and the disappearance of the ultrared matter). \item Long-period (Oort cloud) and short-period (Kuiper belt) comets are statistically identical in their optical color distributions, with no dependency on perihelion distance or other orbital elements. All lack ultrared matter. \item The Jovian Trojan optical colors have an ultrared fraction near 0, inconsistent with a Kuiper belt source. However, while they are apparently not now active, ancient Centaur-like activity driven by their first approach to the Sun with $q \lesssim $ 10 AU could have lead to the loss of ultrared mantle material. \item The same explanation cannot apply to the Neptunian Trojans. They have the same optical color distribution as found in the Jovian Trojans but they are far too distant and cold for any thermal process to reset the surface colors. This Trojan color conundrum presently has no obvious resolution. \end{itemize} | 18 | 8 | 1808.04885 |
1808 | 1808.02883_arXiv.txt | This document is prepared using LaTeX2e\cite{Lamport94} and shows the desired format and appearance of a manuscript prepared for the Proceedings of the SPIE.\footnote{The basic format was developed in 1995 by Rick Hermann (SPIE) and Ken Hanson (Los Alamos National Lab.).} It contains general formatting instructions and hints about how to use LaTeX. The LaTeX source file that produced this document, {\ttfamily article.tex} (Version 3.4), provides a template, used in conjunction with {\ttfamily spie.cls} (Version 3.4). These files are available on the Internet at \url{https://www.overleaf.com}. The font used throughout is the LaTeX default font, Computer Modern Roman, which is equivalent to the Times Roman font available on many systems. | \label{sec:intro} % Begin the Introduction below the Keywords. The manuscript should not have headers, footers, or page numbers. It should be in a one-column format. References are often noted in the text and cited at the end of the paper. \begin{table}[ht] \caption{Fonts sizes to be used for various parts of the manuscript. Table captions should be centered above the table. When the caption is too long to fit on one line, it should be justified to the right and left margins of the body of the text.} \label{tab:fonts} \begin{center} \begin{tabular}{|l|l|} % \hline \rule[-1ex]{0pt}{3.5ex} Article title & 16 pt., bold, centered \\ \hline \rule[-1ex]{0pt}{3.5ex} Author names and affiliations & 12 pt., normal, centered \\ \hline \rule[-1ex]{0pt}{3.5ex} Keywords & 10 pt., normal, left justified \\ \hline \rule[-1ex]{0pt}{3.5ex} Abstract Title & 11 pt., bold, centered \\ \hline \rule[-1ex]{0pt}{3.5ex} Abstract body text & 10 pt., normal, justified \\ \hline \rule[-1ex]{0pt}{3.5ex} Section heading & 11 pt., bold, centered (all caps) \\ \hline \rule[-1ex]{0pt}{3.5ex} Subsection heading & 11 pt., bold, left justified \\ \hline \rule[-1ex]{0pt}{3.5ex} Sub-subsection heading & 10 pt., bold, left justified \\ \hline \rule[-1ex]{0pt}{3.5ex} Normal text & 10 pt., normal, justified \\ \hline \rule[-1ex]{0pt}{3.5ex} Figure and table captions & \, 9 pt., normal \\ \hline \rule[-1ex]{0pt}{3.5ex} Footnote & \, 9 pt., normal \\ \hline \rule[-1ex]{0pt}{3.5ex} Reference Heading & 11 pt., bold, centered \\ \hline \rule[-1ex]{0pt}{3.5ex} Reference Listing & 10 pt., normal, justified \\ \hline \end{tabular} \end{center} \end{table} \begin{table}[ht] \caption{Margins and print area specifications.} \label{tab:Paper Margins} \begin{center} \begin{tabular}{|l|l|l|} \hline \rule[-1ex]{0pt}{3.5ex} Margin & A4 & Letter \\ \hline \rule[-1ex]{0pt}{3.5ex} Top margin & 2.54 cm & 1.0 in. \\ \hline \rule[-1ex]{0pt}{3.5ex} Bottom margin & 4.94 cm & 1.25 in. \\ \hline \rule[-1ex]{0pt}{3.5ex} Left, right margin & 1.925 cm & .875 in. \\ \hline \rule[-1ex]{0pt}{3.5ex} Printable area & 17.15 x 22.23 cm & 6.75 x 8.75 in. \\ \hline \end{tabular} \end{center} \end{table} LaTeX margins are related to the document's paper size. The paper size is by default set to USA letter paper. To format a document for A4 paper, the first line of this LaTeX source file should be changed to \verb|\documentclass[a4paper]{spie}|. Authors are encouraged to follow the principles of sound technical writing, as described in Refs.~\citenum{Alred03} and \citenum{Perelman97}, for example. Many aspects of technical writing are addressed in the {\em AIP Style Manual}, published by the American Institute of Physics. It is available on line at \url{https://publishing.aip.org/authors}. A spelling checker is helpful for finding misspelled words. An author may use this LaTeX source file as a template by substituting his/her own text in each field. This document is not meant to be a complete guide on how to use LaTeX. For that, please see the list of references at \url{http://latex-project.org/guides/} and for an online introduction to LaTeX please see \citenum{Lees-Miller-LaTeX-course-1}. | 18 | 8 | 1808.02883 |
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1808 | 1808.10491_arXiv.txt | In this proceeding, we present studies of instrumental systematic effects for the Simons Obsevatory (SO) that are associated with the detector system and its interaction with the full SO experimental systems. SO will measure the Cosmic Microwave Background (CMB) temperature and polarization anisotropies over a wide range of angular scales in six bands with bandcenters spanning from 27 GHz to 270 GHz. We explore effects including intensity-to-polarization leakage due to coupling optics, bolometer nonlinearity, uncalibrated gain variations of bolometers, and readout crosstalk. We model the level of signal contamination, discuss proposed mitigation schemes, and present instrument requirements to inform the design of SO and future CMB projects. | The cosmic microwave background (CMB) contains a wealth of information in both its temperature and polarization, including information that can help determine the nature of dark energy and dark matter, the mass and number of neutrino species, and if there was a period of inflation shortly after the universe began. The linear polarization of the CMB can be decomposed into even (E-mode) and odd (B-mode) parity components. If inflation occurred, the gravitational waves would have produced both E-modes and B-modes. Since E-modes are also produced by density perturbations, the detection of a B-mode signal would be particularly powerful in giving additional support to inflationary models~\cite{Seljak_1997, Kami_1997}. The amplitude of this primordial B-mode signal is expected to peak on degree-angular scales (multipole moments of $\ell \sim 100$), and its amplitude is quantified by the tensor-to-scalar ratio $r$. There is also a contribution to B-modes at sub-degree scales from the gravitational lensing of E-modes into B-modes from intervening large-scale structure that, in combination with E-modes and temperature anisotropies, holds information about the mass and number of neutrinos and the nature of dark energy and dark matter~\cite{lensing}. However, B-modes are faint, and the primordial B-mode signal has yet to be measured. This has pushed CMB science toward increasingly sensitive experiments where calibration and the mitigation of systematic effects are critical. Further complicating measurements is the contamination of CMB polarization by foregrounds from synchrotron and dust emission. These foreground signals have a different frequency dependence than the CMB polarization, so in principle, these signatures can be removed at the required fidelity if appropriately characterized as a function of frequency. The Simons Observatory (SO) will observe the CMB in both temperature and polarization over a wide range of frequencies (27-270~GHz) and angular scales. SO will field a $\sim$6~m crossed-Dragone, large-aperture telescope (LAT) for observations at small angular scales and several small-aperture ($\sim$0.5~m) telescopes (SATs) for large angular scale measurements. SO plans to push to high sensitivity by deploying $\sim$50 multichroic detector arrays in its initial configuration. This represents more millimeter-wave detectors for the observation of the CMB than have yet been deployed elsewhere and represents a critical step toward next-generation experiments like CMB-S4~\cite{S4}. To fully utilize this sensitivity, the systematic effects must be well-characterized and mitigated through the instrument design. In this paper, we focus on systematic effects originating from the detector system, which include the optical coupling (feedhorns and lenslets), detectors, and signal readout. This paper is part of a series of papers on the systematic and calibration studies for SO~\cite{salatino18,bryan18,gallardo18}. We are combining the detailed results of the full SO systematics and calibration studies into a comprehensive study that will be released in a series of future papers to the community for use in developing future CMB experiments such as CMB-S4. Here we take an in-depth look at several of the most important detector array systematics, and note that further information on detector array systematic effects that are not included in this paper will be included in the full systematics study papers. In Sec.~\ref{sec:analysis_frameworks}, we introduce two analysis frameworks used to model detector systematic effects: a time-domain systematics pipeline and a map-based systematics pipeline. Section~\ref{sec:IP_horn_lenslet} discuses simulations of polarization leakage from feedhorns and lenslets. In Sec.~\ref{sec:nonlinearity} and~\ref{sec:gain_drift}, we discuss the possible spurious polarization induced by two sources of time-varying differential gain between polarization-pair detectors: long-timescale gain drifts and bath temperature fluctuations, respectively. Finally, in Sec.~\ref{sec:crosstalk}, we discuss crosstalk studies for frequency-division readout schemes. | To fully leverage the sensitivity of SO, we must understand and control the level of systematics in the system. Modeling systematics in the design phase of the experiment provides critical feedback to the instrument design to ensure that SO meets its scientific goals. The preliminary studies on detector array systematic effects presented in this work have informed the spacing between pixels, detector parameter selection, the readout layout, and the scan strategy. Studies of the optical coupling to the detector arrays show that the level of polarization leakage is at an acceptable level for the SO pixel sizes on the LAT and SAT with both architectures when beam calibration is included. Going forward, the map-based simulation pipeline developed for this study can be used to study the effect of any generic beam on the power spectra. We have described the relevant models for long-timescale gain drifts that we expect to be sourced by our instrument's interactions with the TES bolometers in the focal plane. These effects are especially problematic for the LAT arrays. In the case of nonlinearity-induced differential gain between polarization-pair bolometers, the leakage of T$\rightarrow$P is not negligible given TES bolometer parameters optimized for background-limited performance and sensitivity on the LAT. However, we expect that, through more extensive future simulations, we can set acceptable upper bounds on these effects to avoid negative impacts on SO science goals by tuning relevant detector parameters. We additionally plan to constrain detector non-linearity through a direct measurement of higher-harmonic response, which will ensure that our TES bolometer arrays can perform as required. Further, we show how bath temperature fluctuations can couple into gain fluctuations of TES bolometers and discuss how this effect can also be controlled with bias power. Studies of both of these effects in the map and spectral domain are advancing toward SO-like detector counts and array configurations. Crosstalk will be an important systematic to consider for the SAT and the LAT readout systems. The SAT will use a CRHWP that will significantly mitigate this effect. It will be important to simulate more realistic focal plane layouts and crosstalk matrix elements based on lab measurements. It will also be important to simulate the combined effect of crosstalk between detectors at different frequencies and spatial locations. The simulations discussed in this work will inform hardware design choices and will incorporate more detail as the SO instrument designs mature. SO is a critical stepping stone for future experiments like CMB-S4. The tools and analyses developed for the SO systematic studies will be publicly released with the SO systematics studies to contribute to the design of future CMB experiments. | 18 | 8 | 1808.10491 |
1808 | 1808.10458_arXiv.txt | In order for Wide-Field Infrared Survey Telescope (\emph{WFIRST}) and other Stage IV dark energy experiments (e.g., Large Synoptic Survey Telescope; LSST, and \emph{Euclid}) to infer cosmological parameters not limited by systematic errors, accurate redshift measurements are needed. This accuracy can be met by using spectroscopic subsamples to calibrate the photometric redshifts for the full sample. In this work we employ the Self Organizing Map (SOM) spectroscopic sampling technique, to find the minimal number of spectra required for the \emph{WFIRST} weak lensing calibration. We use galaxies from the Cosmic Assembly Near-infrared Deep Extragalactic Legacy Survey (CANDELS) to build the LSST+\emph{WFIRST} lensing analog sample of $\rm \sim 36$k objects and train the LSST+\emph{WFIRST} SOM. We find that 26\% of the \emph{WFIRST} lensing sample consists of sources fainter than the \emph{Euclid} depth in the optical, 91\% of which live in color cells already occupied by brighter galaxies. We demonstrate the similarity between faint and bright galaxies as well as the feasibility of redshift measurements at different brightness levels. Our results suggest that the spectroscopic sample acquired for calibration to the \emph{Euclid} depth is sufficient for calibrating the majority of the \emph{WFIRST} color-space. For the spectroscopic sample to fully represent the synthetic color-space of \emph{WFIRST}, we recommend obtaining additional spectroscopy of $\sim 0.2-1.2$k new sources in cells occupied by mostly faint galaxies. We argue that either the small area of the CANDELS fields and the small overall sample size or the large photometric errors might be the reason for no/less bright galaxies mapped to these cells. Acquiring the spectra of these sources will confirm the above findings and will enable the comprehensive calibration of the \emph{WFIRST} color-redshift relation. | Revealing the nature of the dark energy driving cosmic acceleration and testing general relativity on cosmological scales are essential pieces to complete our understanding of modern cosmology and physics. To achieve these goals, the next generation large cosmology surveys will make precision measurements of the expansion history of the Universe as well as the growth rate of large scale structure using various techniques \citep{Spergel2015}. Samples of supernova type Ia (SNIa) constrain cosmological parameters (e.g., the expansion rate of the Universe) by providing measurements of cosmological distances as a function of redshift (e.g., \citealt{Riess1998}, \citealt{Perlmutter1999}). Complementary distance scale measurements can be obtained from baryon acoustic oscillation (BAO) imprints in the power spectrum of cosmic microwave background (CMB) and in the large scale structure (LSS) of galaxies at lower redshifts (e.g., \citealt{Zhan2006}, \citealt{Benitez2009}). Weak gravitational lensing of distant galaxies by the gravitational field of matter inhomogeneities in the large scale structure, or cosmic shear, provides another powerful tool for constraining the power spectrum of dark and luminous matter in the Universe (e.g., \citealt{Blandford1991}, \citealt{Blandford1992}). Weak lensing cosmology requires both redshift estimates and shape measurements of statistical samples of galaxies (e.g., \citealt{Hu1999}). Upcoming stage IV dark energy experiments aimed for 2020s (see \citealt{Albrecht2006}) will improve current measures of distance and cosmic expansion history (with uncertainties $\sim 1-3\%$) as well as matter clustering (with uncertainties $\sim 5-10\%$) to $0.1-0.5\%$ precision, while also extending them to previously unexplored redshift regimes. Careful calibration is required such that the cosmological inferences will not be limited by systematic errors \citep{Spergel2015}. Accurate redshifts are needed for all three techniques mentioned above (SNIa, LSS BAO, and weak lensing tomography). While SNIa and BAO studies usually employ a spectroscopic sample, obtaining spectroscopic redshifts for hundreds of millions to billions of faint galaxies needed for weak lensing analysis is not practical. Therefore, highly accurate photometric redshifts, trained and validated using a training sample of spectroscopic data, are required. Several recent studies (e.g., \citealt{Cunha2012}, \citealt{Newman2015}, \citealt{M15}) have investigated the best spectroscopic sampling strategy in order to train higher-quality, lower-scatter photo-$z$ with less systematic errors for different cosmological surveys. \citet{Carrasco2013} showed that random selection of galaxies to create a spectroscopic training sample is not optimal. Recent work has suggested spatial cross-correlation-based techniques relating the photometric redshifts with a reference spectroscopic sample as a solution (e.g., \citealt{Newman2008}, \citealt{Newman2015}). These techniques also require a large spectroscopic sample. However, their main advantage is that the spectroscopic sample does not need to be representative of different galaxy types (i.e. bright emission line galaxies which are easy spectroscopic targets can be used for calibration). \begin{figure} \centering \includegraphics[trim=0cm 0cm 0cm 0cm, clip,width=0.5\textwidth] {filters.png} \caption{Expected 5$\sigma$ limiting magnitudes of LSST and the \emph{WFIRST} high latitude survey (HLS) filters for galaxies ($\rm r_{1/2}=0.3"$) plotted as solid colored lines. We use deeper or similar depth photometry in the same wavelength range as probed with LSST+\emph{WFIRST} from CANDELS catalogs in five fields to estimate LSST+\emph{WFIRST} photometry. 5$\sigma$ limiting AB magnitude of CANDELS GOODS-S filters for galaxies ($\rm r_{1/2}=0.3"$) are plotted with dotted black lines (CANDELS filters used are listed in Table \ref{tbl:filters}). Euclid's expected riz as well as NIR 5$\sigma$ depths are also over plotted as dashed gray lines for comparison.} \label{fig:filters} \end{figure} A completely data-driven technique of selecting optimal spectroscopic samples to meet the cosmological requirements was introduced by Masters et al. (2015; hereafter M15). This technique uses a machine learning algorithm called the Self Organizing Map (SOM) (Kohonen 1982, 1990) to reduce the multi-dimensional color-space of galaxies defined by a photometric survey to two dimensions (hence maps). This empirical color mapping method allows us to focus our spectroscopic efforts on under-sampled regions of galaxy parameter space. M15 explored different SOM-based targeting strategies and estimated the required spectroscopy for the \emph{Euclid} mission (\citealt{Laureijs2011}). This approach is the basis of a large, ongoing spectroscopic program, the Complete Calibration of the Color-Redshift Relation (C3R2) survey, designed to calibrate the color-redshift relation to the \emph{Euclid} depth (\citealp{Masters2017}, Masters et al. in prep). In this paper, as part of the High Latitude Survey (HLS- \citealt{Dore2018}) science investigation team, we extend the previous analysis of M15 to estimate the additional spectroscopic sample required to meet the Wide-Field Infrared Survey Telescope (WFIRST) cosmological requirements. \emph{WFIRST} is a NASA flagship mission using a 2.4-meter telescope to provide measurements of the expansion history of the universe and growth of structure to better than 1$\%$ (\citealt{Spergel2015}). For weak lensing analysis, the \emph{WFIRST} HLS is currently planing to image $\rm 2227 \ deg^{2}$ in four near infrared (near-IR) bands ($Y$, $J$, $H$, and $F184$) spanning the range from $0.92\!-\!2.00\,\mu\mathrm{m}$ to magnitudes $\rm 25.8\!-\!26.7$ (depending on band), significantly fainter compared to the near-IR depths in the \emph{Euclid} survey ($\sim 24.5$). The near IR filters alone are not sufficient for precise photo-$z$ estimation. Multi-band optical observations need to be combined with \emph{WFIRST} to fulfill the redshift requirements. Such observations will be available through the Hyper Suprime-Cam (HSC) on Subaru or by the Large Synoptic Survey Telescope (LSST). In \S 2 we simulate a data-driven photometric catalog using deep observations from the Cosmic Assembly Near-IR Deep Extragalactic Legacy Survey (CANDELS; \citealt{Grogin2011}, \citealt{Koekemoer2011}) to replicate the LSST+\emph{WFIRST} lensing sample. In \S 3 we briefly review the SOM technique, train the SOM with the LSST+\emph{WFIRST} analog catalog and test for its accuracy in representing the data. We check for the effects of cosmic variance in \S 4 and address the additional spectroscopy needed to meet \emph{WFIRST} cosmology requirement in \S 5. \S 6 summarizes the results of this work and discusses sources of uncertainty. Throughout this paper all magnitudes are expressed in AB system (\citealt{Oke1983}) and we use the concordance $\Lambda$CDM cosmology with $H_{0}=70 \ \rm km \ s^{-1}\ Mpc^{-1}$, $\Omega _{M}= 0.3$, and $\Omega _{\Lambda} = 0.7$. \begin{figure} \centering \includegraphics[trim=0cm 0cm 0cm 0cm, clip,width=0.50\textwidth] {interpol.png} \caption{LSST and \emph{WFIRST} measured photometry for five sample galaxies from CANDELS GOODS-S (squares). Solid lines and circles show the CANDELS broadband filters used for the linear interpolation. Each color represents a sample galaxy.} \label{fig:interpol} \end{figure} | } In this paper, we studied the redshift calibration requirements for \emph{WFIRST} HLS weak lensing analysis to meet the stage IV dark energy desired accuracy. We adopted the methodology introduced by M15, which calibrates the color-redshift relation using SOMs. We imitated the LSST+\emph{WFIRST} lensing sample using optical and near-IR data from the five CANDELS fields and trained a SOM with successive colors of galaxies in the LSST+\emph{WFIRST} filter set. The smoothness of redshift distribution on a SOM trained by colors, illustrates the color-redshift relation and makes the SOM an optimal source for spectroscopic target selection. Based on Monte Carlo simulations in M15, and given the estimated average redshift uncertainties in our SOM cells, a tomographic bin containing $\sim 200$ SOM cells, would be sufficient to reach $\Delta \langle z \rangle / (1+\langle z \rangle) < 0.002 $. For the technique to be efficient, most SOM cells need to have at least one spectroscopic object mapped to them for calibration. This is equivalent to $\sim 5 $k total spectroscopic redshifts to calibrate the \emph{WFIRST} SOM. However, in addition to the already existing spectroscopic observations in the CANDELS fields (covering $57\%$ of the SOM cells) the C3R2 survey is filling the color-space of \emph{Euclid} galaxies with spectroscopic observations. We showed that $\sim 26\%$ of the \emph{WFIRST} lensing sample consists of sources fainter than the \emph{Euclid} depth in the optical, $91\%$ of which live in color cells also occupied by brighter (\emph{Euclid}-depth) sources. We demonstrated the similarity between the fainter and brighter subsamples in same cells as well as the feasibility of measuring the redshifts of fainter objects to the accuracy needed using the SOM color-redshift relation. Since the $\sim 4\%$ of cells which have only fainter objects associated to them, might be due to small sample size in the CANDELS fields as well as larger photometric errors in the fainter sample, we recommend extra spectroscopy for these cells to calibrate the color-redshift relation on \emph{WFIRST} SOM throughly. We recommend $\sim 0.2-1.2 $k new spectra, which will cover the cells with only faint objects as well as those with large redshift dispersions. It is crucial to note that, having most of the calibration already in place by the C3R2 \emph{Euclid} effort does not imply similarity between \emph{WFIRST} weak lensing cosmology and the less deep surveys, as the lensing sample size will increase significantly (see Figure \ref{fig:redshifts}). \begin{figure} \centering \includegraphics[trim=0cm 0cm 0cm 0cm, clip,width=0.49 \textwidth] {WFIRST_EUCLID_zdist.png} \caption{\emph{WFIRST} will significantly increase the lensing sample size. Normalized redshift distribution of galaxies in the LSST+\emph{WFIRST} analog sample is shown as the blue histogram and the fraction of galaxies with $riz<25$ as would be found by \emph{Euclid} is over plotted in purple.} \label{fig:redshifts} \end{figure} In our analysis, we have used an interpolation technique to estimate photometry in LSST+\emph{WFIRST} filter sets based on available photometry in different filter sets. The method is tested extensively and is the easiest logical way to reproduce statistically correct distributions of photometries and colors of galaxies. Future works using improved techniques will test for the robustness of the interpolation on an object by object basis. One weakness of this work is the small sample size used for training the SOM. As discussed in the paper, CANDELS data were the best available option due to comparable depth of observations to those expected from the \emph{WFIRST} HLS. CANDELS observations being done in five well-separated fields in the sky should mitigate the effect of cosmic variance. We found the color-space of our \emph{WFIRST} lensing sample to be representative of a sample trained for Euclid (19 times larger in area), which suggests that the effect of cosmic variance in the SOM calibration technique should be minimal. However, very large area simulations are needed to enable a more quantitative investigation of this effect on the findings presented in this work (Capak et al. in prep.). We will revisit our forecasts once \emph{WFIRST} observations are available, and retrain a SOM with actual observations for photometric redshift calibration as well as selection of weak lensing tomographic bins. | 18 | 8 | 1808.10458 |
1808 | 1808.09643_arXiv.txt | Recently, the nonperturbative quantization scheme of loop quantum gravity has been extended to the Brans-Dicke theory and the corresponding loop quantum Brans-Dicke cosmology has been derived, which provides an essential platform to explore inflationary models in this framework. In this paper, we consider two inflation models, the Starobinsky and $\alpha$-attractor inflation whose cosmological predictions are in excellent agreement with Planck data, and study systematically their pre-inflationary dynamics as well as the slow-roll inflation. We show that for both models, the background evolution of a flat Friedmann-Lema\^{i}tre-Robertson-Walker universe in general can be divided into three different phases: the pre-inflationary quantum phase, quantum-to-classical transition, and the slow-roll inflation. The pre-inflationary dynamics are dominated by the quantum geometry effects of loop quantum Brans-Dicke cosmology and the corresponding Universe could be either initially expanding or contracting, depending on the initial velocity of inflaton field. It is shown that the detailed evolution of pre-inflationary quantum phase also depend on specific inflation models. After the pre-inflationary quantum phase, the universe gradually evolves into the slow-roll inflation with some of initial conditions for Starobinsky and $\alpha$-attractor potentials. In addition, to be consistent with observational data, we also find the restricted parameter space of initial conditions that could produce at least $60$ $e$-folds during the slow-roll inflation. | \renewcommand{\theequation}{1.\arabic{equation}}\setcounter{equation}{0} The paradigm of cosmic inflation provides perhaps the most compelling picture of the universe at the early stages of its history. It has achieved remarkable successes not only in solving several problems of the standard big bang cosmology, but most importantly in predicting the primordial perturbations spectra whose evolutions explain both the formation of the large scale structure of the universe and the tiny anisotropies in the cosmic microwave background (CMB) \cite{guth_inflationary_1981, sato_firstorder_1981, starobinsky_new_1980}. All these predictions are all matched to cosmological observational data with high precisions \cite{komatsu_sevenyear_2011, planckcollaboration_planck_2018, planck_collaboration_planck_2014-1, planck_collaboration_planck_2015-4}. In general, there are a lot of approaches to realize the inflation that originate from very different background physics. A common trait of many inflationary models is that they involve scalar degrees of freedom with a self-interacting potential that gives rise to a slow-roll phase during which the energy density of the matter field remains nearly constant and the spacetime behaves like a quasi-de Sitter spacetime. Recently, cosmological and astrophysical data show that for single field inflationary models, predictions of the Starobinsky and $\alpha$-attractor inflation with a small value of $\alpha$ are favored over others \cite{planckcollaboration_planck_2018}. The Starobinsky inflation is based on the account of $R^2$-term as the correction in the Einstein equations \cite{starobinsky_new_1980}, which emerges in the Planck epoch and plays a fundamental role in the high curvature limit, when the early-time acceleration takes place. The $\alpha$-attractor inflation can be in general considered as extensions of the Starobinsky inflation from $\alpha=1$ to $\alpha \neq 1$ \cite{kallosh_superconformal_2013, ferrara_minimal_2013}. Such theories are conformally equivalent to a scalar-tensor theory in the Einstein frame, where the inflaton drives the expansion in a quasi-de Sitter space-time and slowly moves to the end of inflation \cite{tsujikawa_planck_2013, felice_theories_2010}. As expected, their perfect agreement with Planck data has renewed interest in these models. However, inflationary theory itself is very sensitive to physics at Planck scales, due to the fact that the energy scale of inflation may not be far from that of quantum gravity \cite{martin_transplanckian_2001, zhu_uniform}. Because of this, the underlying quantum field theory and classical general relativity (GR) on classical spacetime becomes unreliable for a large class of inflationary models. This is also known as the ``trans-Planckian issues" of the inflationary theory and its implications on primordial perturbation spectra have also been studied in some concrete quantum theories of gravity, for example see the discussions in Horava-Lifshitz gravity \cite{zhu_hovara}. In addition, insisting on the use of classical GR to describe the inflationary process will inevitably lead to an initial singularity \cite{borde_eternal_1994, borde_inflationary_2003}. All these issues are closely related to the regime where classical GR is known to break down, and one expects a quantum theory of gravity will provide a completed description of inflation as well as its pre-inflationary dynamics. To address these issues, Loop quantum cosmology (LQC) provides a natural framework, in which the standard inflationary scenarios can be extended from the onset of the slow-roll inflation back to the Planck era in a self-consistent way. In such a picture, the big bang singularity is replaced by a finite nonzero universe, the quantum bounce, which eventually evolves to the desired slow-roll inflation with very high possibilities \cite{singh_nonsingular_2006, zhang_inflationary_2007, chen_loop_2015, zhu_preinflationary_2017a, agullo_quantum_2012}. Such remarkable features of the quantum bounce have attracted a great deal of attentions lately, in which the universe dominated by a scalar field for different scalar field potentials have been widely explored \cite{ye_loop_2018, zhu_preinflationary_2017a, chen_loop_2015, shahalam_preinflationary_2018, li_qualitative_2018, li_cosmological_2018, shahalam_preinflationary_2017, sharma_preinflationary_2018, agullo_preinflationary_2013, agullo_detailed_2015}. In addition, the primordial perturbations spectra with loop quantum corrections and effects of quantum bounce and their observational constraints have also been discussed extensively (see \cite{zhu_primordial_2018a, zhu_preinflationary_2017a, zhu_universal_2017b, agullo_preinflationary_2013, agullo_loop_2015, agullo_detailed_2015, ashtekar_quantum_2017, zhu_uniform_lqc} and references therein). Since the Starobinsky and $\alpha$-attractor inflation models are favored by Planck data, an essential question arising in the framework of LQC is to check if the slow-roll inflation for the Starobinsky and $\alpha$-attractor models can still occur after the quantum bounce at the Planck era. In fact, both the dynamics of background and cosmological perturbations of Starobinsky inflation model have been already studied in the framework of LQC of GR \cite{bonga_phenomenological_2016, zhu_preinflationary_2017a}. One of main conclusions of these studies is that following the quantum bounce in LQC of GR, a desired slow-roll inflation phase is almost inevitable and the imprints of quantum bounce on primordial perturbation spectra and non-Gaussianities can be well within observational constraints \cite{zhu_universal_2017b, zhu_preinflationary_2017a, zhu_primordial_2018a, wu_nonadiabatic_2018}. However, as mentioned in \cite{zhu_preinflationary_2017a}, most of the above mentioned studies about Starobinsky inflation are limited to the effective dynamics obtained from the loop quantization in the Einstein frame. In this frame, the original theory of Starobinsky inflation and its extensions ($\alpha$-attractors) in the Jordan frame have been transformed into the Einstein frame by using a conformal transformation, so that the slow-roll inflation can be driven by a scalar field with the specific potentials in the framework of GR. Classically this is correct because the descriptions of the slow-roll inflation in both frames are equivalent. However, whether this is also true or not in LQC is still an open question. According to \cite{artymowski_comparison_2013}, in general the Einstein and Jordan frames are non longer equivalent at the quantum level. Thus, it is interesting to explore the inflation directly with the effective dynamics obtained from the loop quantization directly in the Jordan frame, based on the quantization proposed in \cite{zhang_loop_2012, zhang_loop_2013a,zhang_extension_2011, zhang_loop_2011,zhang_nonperturbative_2011}. In general, both the theories of Starobinsky and $\alpha$-attractor inflation can be casted into the form of specific types of Brans-Dicke (BD) theory in the Jordan frame \cite{tsujikawa_planck_2013, felice_theories_2010}. Recently, the nonperturbative quantization scheme of LQG has been successfully extended to the BD theory \cite{zhang_loop_2012, zhang_loop_2013a}. The corresponding effective equations of cosmological model for loop quantum BD cosmology have been derived \cite{zhang_loop_2013a} based on the loop quantization procedure in the Jordan frame. These effective equations thus provide a very essential platform to study the Starobinbsky and $\alpha$-attractor inflation in the framework of loop quantum BD cosmology. To this purpose, in this paper we study the Starobinsky and $\alpha$-attractor inflation as well as their pre-inflationary dynamics in the framework of loop quantum BD cosmology. Because the quantization in different frames may give different results, we expect loop quantum BD cosmology may provide some distinguishing description about the background evolutions of Starobinsky and $\alpha$-attractor inflation from those in LQC of GR in the Einstein frame. Since most of studies and results about Starobinsky and $\alpha$-attractor inflation are considered in the Einstein frame, here in order to compare our results with theirs, in this paper we shall focus on quantities of background evolution in loop quantum BD cosmology (in Jordan frame) by writing them in terms of those in LQC of GR (in Einstein frame). With this strategy, we show that the evolution of the background in general can be divided into three phases: {\em the pre-inflationary quantum phase, quantum-to-classical transition, and the slow-roll inflation}. For pre-inflationary quantum phase, we shall observe that the Universe starts at a finite non-zero Universe which could be either contracting or expanding depending on the initial velocity of scalar field $\chi$. For slow-roll inflation, we also show the parameter space that could leads to at least $60$ $e$-folds during the slow-roll inflation. This paper is origanized as follows. In Sec. II, we give an introduction of the classical dynamics of slow-roll inflation in BD theory and in particular focus on the Starobinsky and $\alpha$-attarctor inflation as well as their observational constraints. In Sec. III, we present the effective equations about the background evolution in loop quantum Brans-Dicke cosmology and then transform them in terms of quantities in the Einstein frame. Based on these effective equations, in Sec. IV we turn to study the background evolution by using numerical calculations in details for both Starobinsky and $\alpha$-attractor inflation. Our main conclusions and discussions are presented in Sec. V. | \renewcommand{\theequation}{5.\arabic{equation}}\setcounter{equation}{0} In this paper we have provided a detailed numerical study of the Starobinsky and $\alpha$-attractor inflation as well as their pre-inflationary dynamics in the framework of loop quantum BD cosmology. We show that for both the Starobinsky and $\alpha$-attractor inflation, the evolution of the background Universe can be in general divided into three different phases: {\em pre-inflationary quantum phase, quantum-to-classical transition, and slow-roll inflation}. During the pre-inflationary quantum phase, the background evolution is dominated by the quantum geometry effects of loop quantum BD cosmology. Unlike the background evolution in LQC of GR where the pre-inflationary dynamics represents an expanding Universe started at the quantum bounce \cite{zhu_preinflationary_2017a}, the pre-inflationary dynamics is very complicated and depends on initial conditions and specific models. Generally, the Universe is initially expanding if the initial velocity of the scalar field $\chi$ is positive ($\dot \chi_0>0$) and is contracting if it is negative ($\dot \chi_0 <0$). For Starobinky inflaton($\beta=3$) and $\alpha$-attractor inflation with $\beta=1$, the initial expanding Universe shall collapse to a contracting phase before it evolves into the final expanding phase through the quantum bounce, while initial contracting Universe directly connects to the expanding phase through the quantum bounce. For $\alpha$-attractor inflation with $\beta=5, \; 10, \; 20$, we show that the quantum bounce does not exist after the initial time $\hat t_0$ for the initial expanding Universe. Whether a quantum bounce would appear before the chosen initial time $\hat t_0$ deserve further investigating. This issue concerns whether the effective equations are still valid for extremely high energy near to the classical singularity and thus is out of the scope of this paper. For initial contracting Universe, the evolution of the background is almost the same as that in models of Starobinsky inflation ($\beta=3$) and $\alpha$-attractor inflation with $\beta=1$. After the pre-inflationary quantum phase, the universe gradually evolves into the expanding Universe. For some of initial conditions in the parameter space, we show that the slow-roll inflation for both the Starobinsky and $\alpha$-attractor models are produced. In addition, to be consistent with observational data, we also derive the range of initial conditions that could produce at least $60$ $e$-folds during the slow-roll inflation. | 18 | 8 | 1808.09643 |
1808 | 1808.10344_arXiv.txt | Several hydrodynamical instabilities potentially operating in disks around young stars have been discussed lately that all depend on the thermodynamic stratification of the disk and on the local rate of thermal relaxation. In this paper, we map constraints on the spatial extents of hydrodynamically unstable regions for the {\em Vertical Shear Instability} (VSI), the {\em Convective OverStability} (COS), and the amplification of vortices via the {\em Subcritical Baroclinic Instability} (SBI) onto models of viscous accretion disks. % We use 1+1 dimensional, steady state accretion disk models, which include the effects of stellar irradiation, viscous heating and flux-limited radiative transfer. This allows for a self-consistent determination of the local radial and vertical stratification and thermal relaxation rate in a disk, in dependence of the chosen parameters like stellar mass, disk mass and the viscosity parameter $\alpha$. Our results imply that passive disks with $\alpha\lesssim 10^{-4}$ do only have VSI and COS susceptible zones at radial distances $\gtrsim 10\, \text{au}$ and about one pressure scale height above the midplane. Vortex amplification via SBI nevertheless operates even for this low viscous heating rate between $\sim 0.3\, \text{au}$ and $\sim 50\, \text{au}$. For $\alpha \gtrsim 10^{-4}$, VSI and COS become active down to radial distances of $\sim 1\, \mathrm{au}$. Here, despite relatively long thermal relaxation times, VSI can operate even close to the midplane, because of the vertically adiabatic stratification of the viscously heated disk. The growth of all considered instabilities (VSI, COS, and SBI) is favored for lower stellar masses and larger disk masses and $\alpha$-values. \CORR{We conclude that if any mechanism is able to create at least low $\alpha$-stresses ($\sim 10^{-4}$) that lead to viscous heating of the disk, hydrodynamical instabilities are likely to operate in significant parts of the planet forming zones in disks around young stars.} Thus these instabilities are possible candidates to explain the rings and vortices observed with ALMA and VLT. | \label{sec:intro} Angular momentum transport and the associated accretion process in protoplanetary disks are either driven by winds \citep{Wardle_1997, Pudritz_2007, Koenigl_2010, Bai_2013} or by magnetic and hydrodynamic turbulence \citep{Luest_1952, Shakura_1973, Balbus_Hawley_1991}. One of the considerable physical processes, causing outward transport of angular momentum is the Magnetorotational Instability (MRI), which requires a sufficiently ionised shear flow in addition to weak magnetic fields. This linear instability works well in accretion disk of high temperature around black holes or neutron stars. Large parts of protoplanetary disks however have low ionisation rates and gas densities outside $\sim \SI{0.3}{\AU}$, thus non-ideal MHD effects, namely resistivity and ambipolar diffusion \citep{Lesur_2014, Gressel_2015}, largely hamper the MRI and thus open a venue for hydrodynamic instabilities, as explored by \cite{Lyra_2011}. These hydrodynamical mechanisms include the Subcritical Baroclinic Instability (SBI) \citep{Klahr_Bodenheimer_2003, Petersen1_2007, Petersen2_2007, Lesur_Papaloizou_2010}, the Convective Overstability (COS) \citep{Klahr_Hubbard_2014}, which can be interpreted to be the linear phase of the SBI-mechanism \citep{Lyra_2014}, and the Vertical Shear Instability (VSI) \citep{Urpin_1998, Urpin_2003, Nelson_2012, Lin_2015}, which is the protoplanetary disk equivalent of the Goldreich-Schubert-Fricke Instability \citep{Goldreich_Schubert_1967, Fricke_1968} in stars. The radial and vertical stratification of the disk in temperature and density and the thermal relaxation timescale decide on whether these instabilities can exist or not. For infinite cooling times the stability constraints are given by the standard Solberg-H\o iland criteria \citep{Ruediger_2002}. \cite{Malygin_2017} have investigated detailed models of the radiative properties of a simple, non accreting powerlaw disk profile, identifying necessary conditions for the onset of instability by mapping where the infinite cooling time condition is sufficiently violated. In this paper, we replace the powerlaw disk models with a self-consistent 1+1D accretion disk model that allows to determine the non-trivial temperature and density stratification of the gas as a result of gas accretion \citep{Meyer_1982, Bell_1997} and stellar irradiation \citep{DAlessio_1998}. Whereas the surface temperature of a disk can usually be nicely approximated by a power-law, the midplane temperature can have a more complicated structure with varying gradients, reflecting the local optical depth and the rate of viscous heating \citep{Bell_1997,DAlessio_1998}. \ADD{Gas accretion is assumed to be the result of turbulent viscosity, determined by the free disk parameter $\alpha$ \citep{Shakura_1973}, which for our model determines the amount of thermal energy that is released insides of the disk. The goal of this investigation is therefore not to trace back the origin of $\alpha$ to one of the instabilities we investigate, but to use its pre-defined value to set-up different disk structures in order to study under which conditions and in which kinds of accretion disks the investigated instabilities might be able to grow and thus produce additional substructure or turbulence on top of the back-ground state. We are therefore probing $\alpha$ values between $10^{-5}$ (almost quiescent) and $10^{-2}$ (highly turbulent), which can be thought to be the results of either another instability like the MRI (which would probably quench the instabilities we are aiming to investigate) or be produced by magnetically driven disk winds or Hall MHD, leading to the dissipation of kinetic energy.} The knowledge of the physical conditions inside of a protoplanetary disk makes it possible to determine where the necessary criteria for instability are met and how fast the corresponding linear perturbations grow with time. This is crucial to understand the nature of angular momentum transport in circumstellar disks and to set-up simulations of hydrodynamic instabilities in realistically modeled physical environments. The weak hydrodynamical instabilities are also of special interest for planet formation theory, because even if their contribution to angular momentum transport might be small, they drive the formation of non laminar flow features like zonal-flows and vortices. Such pressure maxima are able to accumulate the inwards drifting dust particles and could therefore be the birthplaces of planetesimals and planets \citep{Barge_1995, Klahr_2006}. Such structures are observed lately in circumstellar disks by ALMA and VLT \citep{Marel_2013, Carrasco_2016}. For a recent review on the role of non-laminar flow features on planetesimal formation we refer to \cite {Klahr2018}. In the following section we give an overview of the basic physics of the investigated instabilities, their analytical growth rates, and the concepts of thermal relaxation used in the scope of this work. The 1+1D disk model as well as the used opacity model are described in \autoref{sec:models}. The general influence of the disk structure on stability, as well as the spatial distributions and growth rates of the introduced mechanisms are presented in \autoref{sec:results}. Our stability maps, shown in \autoref{subsec:stability} sum up the gained knowledge of the distribution of the susceptible regions for the investigated instabilities and parameter set. One example of such a map is shown in \autoref{map}, for a disk model with input parameters $M_{\mathrm{disk}}=\SI{0.1}{\solmass}$, a moderate local viscous heating of $\alpha=10^{-3}$ and a central star of $M_*=\SI{1}{\solmass}$. We finally summarise and conclude in \autoref{sec:conclusion}. \begin{figure}[ht] \centering \includegraphics[width=0.455\textwidth]{Map.pdf} \caption{Stability Map for a Solar Nebula like disk around a solar mass star with $M_{\mathrm{disk}}=\SI{0.1}{\solmass}$ and a moderate local viscous heating of $\alpha=10^{-3}$. The black lines indicate 1, 2 and 3 pressure scale heights respectively. {\bf \color{seagreen} Vertical Shear Instability (VSI - green lines slanted to the right)} can occur at larger radii, where cooling is efficient or alternatively in the inner more optical thick parts of the nebula, which is vertically adiabatic, indicated by the possible occurrence of {\bf \color{maroon} Vertical Convective Motions (VCI - vertical red lines)}. {\bf \color{dodgerblue} Convective Overstability (COS - horizontal blue lines)} will occur in the irradiation dominated outer parts only at a height of 1 pressure scale height and above, because the entropy gradient flips sign only some height above the midplane. In the inner parts of the nebula, where viscous heating is important and especially in the opacity transition zones (ice line \SIrange{3}{6}{\AU} and silicates evaporation zones \SIrange{0.4}{1.5}{\AU}) the entropy gradient is negative down to the midplane, forming a sweet-spot for COS. The {\color{gray}\bf Convective Amplification of large scale vortices (SBI - Grey Shaded region)} can occur through out most of the disk (\SIrange{0.3}{50}{\AU}), because the surface density structure is shallower than the midplane density, thus radial buoyancy (negative gradient of vertical integrated entropy) can much easier be established.} \label{map} \end{figure} | \label{sec:conclusion} In this paper, we investigate the stability of active protoplanetary disks by use of 1+1D steady state accretion disk models including stellar irradiation. This allows for the treatment of flux-limited radiative transfer, caused by viscous heating and makes it possible to apply detailed models of the local rate of thermal relaxation, thus making it possible for the first time to the authors' knowledge, to spatially map the growth rates of the COS, the SBI and the VSI to the radial-vertical plane of a realistically stratified circumstellar disk. We found that we can reproduce the radial temperature profile in full time-dependent 3D axissymmetric radiation hydro simulations as performed by \cite{Bitsch2015} for their model parameters. It is shown that even a quiescent disk, with an extremely low viscosity parameter of $\alpha=10^{-5}$ becomes unstable to COS and VSI at radii $\gtrsim \SI{10}{\AU}$ and for the COS even closer to the star at heights of $\sim \SI{1}{\scaleheight}$ above the midplane. The turbulent structures, resulting of such instabilities were shown to grow to large scale vortices by various authors \citep{Meheut_2012, Lyra_2014, Manger_2018}. A closer look was therefore also taken upon the growth rates and susceptible regions of these structures by consideration of the SBI mechanism. Our results show, that the vertically integrated, radial stratification of the disk allows for positive growth rates over almost the whole radial extent of the disk, between $\sim \SIrange{0.3}{50}{\AU}$ at timescales of $\sim \SI{1000}{\orb}$ even if $\alpha$ has low values of $10^{-5}-10^{-4}$. Vortices which evolve at large radii, due to the perturbations caused by COS and VSI might thus be able to migrate towards the central star while being constantly forced by the SBI contributing to their longevity \citep{Lesur_Papaloizou_2010, Paardekooper_2010}. These findings also indicate that the SBI-mechanism could be a controlling mechanism for the frequently observed vortices in protoplanetary disks \citep{Carrasco_2016,Marel_2013}. Inner regions of the disk, which reach $\alpha$-values of $10^{-4}$, are already subject of strong enough viscous heating and evolve into a radially and vertically buoyant structure, as was shown in \autoref{subsec:cos} and \autoref{subsec:vci}. We have shown that these buoyantly unstable zones arise wherever the disk becomes optically thick enough to allow for the existence of strong vertical temperature gradients which means that they are larger in disks with a high mass ($\gtrsim \SI{0.01}{\solmass}$ for a $\SI{1}{\solmass}$ star). The resulting radial buoyancy, in combination with thermal relaxation is able to drive COS at maximal growth rates of $\sim \SI{e-3}{\om}$. COS is generally favored by massive disks around low mass stars and high $\alpha$-viscosity. \CORR{The opacity structure and thus the relaxation time criterion determine the fastest growing modes, which means that short wavelength perturbations (large $k$) grow best in optically thick regions and long wavelength perturbations grow better in optically thinner regions.} VSI depends even stronger on radiative cooling and thus requires the disk to be at least buoyantly neutral, which impedes repelling forces on vertical perturbations. The internal heat production, caused by $\alpha$ values of $10^{-4}$, is sufficient to provide such a disk structure in the denser interior zones and therefore makes it possible for the VSI to arise, even if the necessary criterion for cooling by \cite{Lin_2015} is not fulfilled. The finding that VSI can produce an $\alpha = 10^{-4}-10^{-3}$ \citep{Nelson_2012, Stoll_2014,Manger_2018} enables the possibility that the VSI can maintain the thermal structure of the disk that it needs to operate. What needs to be shown, is at what height in the disk, the thermal energy will actually be released. Two distinct regions of VSI growth exist at small radii between $\sim \SIrange{0.3}{1.1}{\AU}$ and $\sim \SIrange{3}{3.2}{\AU}$ (for model parameters of $M_{*}=\SI{1}{\solmass}$, $M_{\mathrm{disk}}=\SI{0.1}{\solmass}$, and $\alpha=10^{-3}$), which coincide with the regions unstable to vertical buoyancy. Larger disk masses and $\alpha$-parameters enhance viscous heating and therefore also increase the spatial extents of these inner susceptible zones. Fast cooling is required to allow for VSI far away from the central object, where the vertical stratification is stabilised by weak stellar irradiation and the absence of viscous heating. Thermal relaxation in these parts of the disk is limited by either collisional or radiative timescales. We have shown, that low mass disks ($M\lesssim \SI{0.01}{\solmass}$) have too small gas densities at radii $\gtrsim \SI{10}{\AU}$ to allow for frequent enough collisions between emitters and carriers of thermal energy to cool the disk sufficiently fast. The VSI is therefore strongly hampered in such setups. The gas far away from low mass stars also has low rates of thermal relaxation, since the low temperatures of gas and dust particles do not allow for efficient emittance of energy via black body radiation. We have therefore shown that VSI at large distances from the star is strongly present for large $M_{*}$ and $M_{\mathrm{disk}}$ whereas its presence at smaller distances is determined by the vertically polytropic structure, and therefore favored by small $M_{*}$, large $\alpha$ and $M_{\mathrm{disk}}$. The growth rate of the instability increases vertically from $\sim \SIrange{e-4}{e-2}{\om}$. Future work should deal with more recent opacity models, including the evolution of the dust component of the disk, and a variable chemical composition of the disk, having influence on the adiabatic coefficient $\gamma$. Especially the latter one probably has an important influence on the strength of buoyancy driven instabilities. Passive disks with an assumed surface density profile inspired by the so called minimum mass solar nebula \citep{Chiang1997}, i.e. $\beta_\Sigma = -1.5$ and radial temperature gradient determined only by irradiation will have a radial increasing entropy structure and therefore will be not be the subject of SBI. But note that this steep density profile was extrapolated from the ``solid'' mass distribution in our solar system and not from the gas distribution around the young sun, which should be completely different in the course of dust growth and pebble drift \citep{Birnstiel2012}. Modeling of an actively accreting disk, be it as simple as in our model, or a more elaborate 1+1D irradiation model as in \cite{Dalessio2006} or even in a full 2D hydrodynamic model gives a much shallower surface density profile than $\beta_\Sigma = -1.5$, in range of the values derived from observations \citep{Andrews2010} of $\beta_\Sigma = -1.1$ to $-0.4$. This value predicts a wide region in a disk to have a negative, buoyantly unstable gradient in vertical integrated entropy, i.e. the condition for convective amplification of vortices (SBI). Our model is in general agreement with the more complicated models of disk structure \citep{Dalessio2006,Bitsch2015} if one considers about the first two pressure scale heights above the midplane. Our predictions for stability/instability above that region is less reliable. Here we suggest additional work to investigate cooling rates and entropy structures in these dilute regions, but remember that those regions might also be influenced by magnetic fields because of sufficient ionisation state \cite{Dzyurkevich_2013}, yet being hampered by ambipolar diffusion. Typically those would be the regions where a wind is being launched from the disk in an interplay of photo-evaporation and magnetic fields \citep{Pudritz_2007,Koenigl_2010}. We find that hydrodynamical instabilities can exist in large portions of protoplanetary disks provided that at least a fraction of the released heat gets deposited within 1 or 2 pressure scale heights around the midplane. Even disks with very low accretion rates have a radial temperature stratification, which renders them unstable to SBI and partially unstable to VSI and COS, especially at radii $\gtrsim \SI{10}{\AU}$. The resulting turbulence was shown to be able to create vortices by various other authors \citep{Manger_2018,LyraKlahr_2012,Lyra_2014}. Those structures would be able to migrate through large parts of the disk \citep{Paardekooper_2010} and survive due to their amplification via the SBI-mechanism \citep{Klahr_Bodenheimer_2003,Petersen1_2007,Petersen2_2007,Lesur_Papaloizou_2010}. If the resulting $\alpha$-values are low to moderate ($\alpha=10^{-4}-10^{-3}$), as suggested by \cite{2013ApJ...765..115R} also regions closer to the star become vertically and radially buoyant and therefore susceptible for COS and VSI. Thus, the largest caveat in our work is, that even if sufficient $\alpha$ values are measured in simulations of SBI and VSI \citep{2013ApJ...765..115R, Nelson_2012,Stoll_2014} it is currently not known where the kinetic energy resulting from the release of potential energy in the accretion disk is deposited. If at least a part of this energy is deposited close to the midplane of the disk, then hydrodynamical turbulence has a good chance to operate in large parts of the planet forming regions of protoplanetary disks, where non-ideal MHD effects damp otherwise dominant magnetic effects sufficiently \citep{LyraKlahr_2012}. With its well known properties of forming vortices \citep{2013ApJ...765..115R,Manger_2018} and zonal flows, hydrodynamic instability can be a major agent in forming planetesimals and thus determine the properties of planetary systems \citep{Klahr2018}. \clearpage | 18 | 8 | 1808.10344 |
1808 | 1808.07835_arXiv.txt | {The nearest known binary brown dwarf \object{WISE J104915.57-531906.1AB} (LUH 16) is a well-studied benchmark for our understanding of substellar objects. Previously published astrometry { of LUH 16 obtained with FORS2 on the Very Large Telescope} was affected by errors that limited its use in combination with other datasets, thereby hampering the determination of its accurate orbital parameters and masses. We improve upon the calibration and analysis of the FORS2 astrometry with the help of Gaia DR2 to generate a high-precision dataset that can be combined with present and future LUH 16 astrometry. We demonstrate its use by combining it with available measurements from the Hubble Space Telescope (HST) and Gemini/GeMS and deriving updated orbital and mass parameters. Using Gaia DR2 as astrometric reference field, we derived the absolute proper motion and updated the absolute parallax of the binary to 501.557$\pm$0.082~mas. We refined the individual dynamical masses of LUH 16 to $33.5\pm 0.3\, M_{Jup}$ (component A) and $28.6\pm 0.3\,M_{Jup}$ (component B), which corresponds to a relative precision of $\sim$1\% and is three to four times more precise than previous estimates. { We found that these masses show a weak dependence on one datapoint extracted from a photographic plate from 1984. The exact determination of a residual mass bias, if any, will be possible when more high-precision data can be incorporated in the analysis.} } | After its discovery by \citet{Luhman}, the binary brown dwarf WISE J104915.57-531906.1 (Luhman 16, hereafter LUH 16) located at 2 pc from the Sun has been observed extensively. The astrometric follow-up by \citet{Boffin} used observations with the FORS2 camera on the Very Large Telescope (VLT) between April and June of 2013 to refine the parallax value to $2.020\pm 0.019$ pc and to claim indications for the presence of a massive substellar companion around one of the binary components. Including additional { epochs obtained in 2014 with FORS2,} \citep{2015MNRAS} derived the relative (500.23 mas) and absolute parallax $\varpi_\mathrm{abs}=500.51 \pm 0.11$~mas of this system, obtained an upper mass limit for a potential third body of 2$M_{\rm Jup}$, and determined a mass ratio $q=0.78 \pm 0.10$ for the LUH 16 binary. After publication of Gaia Data Release 1 \citep[DR1][]{GaiaCollaboration:2016aa, DR1_coll, Lindegren:2016aa}, we updated the FORS2 distortion correction, which led to the updated values of $501.139$ mas and $\varpi_\mathrm{abs}=501.419 \pm 0.11$~mas for the relative and absolute parallax, respectively \citep{gaia}. Using Hubble Space Telescope (HST) observations in 2014--2016, \citet{Bedin} derived a relative $501.118 $~mas and absolute parallax $501.398 \pm 0.093$~mas for LUH 16. Because of a longer observation time span, they were able to obtain estimates of the orbital elements of the system, in particular, $31.3\pm 7.9$ years for the orbital period and $0.463 \pm 0.064$ for eccentricity. \citet{Bedin} found inconsistencies between their HST astrometry and the \citet{2015MNRAS} astrometry at the level of 10--20~mas and therefore did not use FORS2 measurements for the orbit fitting. \citet{Garcia} performed an independent analysis of the FORS2 observations obtained in 2013--2014 and added five more epochs in 2015, using Gaia DR1 for the transformation from CCD positions to the International celestial reference frame (ICRF). They added relative astrometry from the Gemini South Multiconjugate Adaptive Optics System (GeMS, \citealt{Ammons}), {ESO photographic plates,} { CRIRES radial velocity,} and archival observations from the Deep Near-Infrared Survey of the Southern Sky (DENIS). The HST astrometry was not published at that time and was not used. Here we report new results from the combination of all available datasets, which increases the time span of high-precision astrometry from $\sim$2 years covered by the \citet{Bedin} study to 3.5 years. In comparison with our first study \citep{2015MNRAS}, we present improved astrometric calibrations of the FORS2 data. We then update the orbital parameters and dynamical masses of the binary on the basis of a combination of FORS2, HST, GeMS, {ESO archive,} Gaia DR1, and Gaia DR2 \citep{DR2, 2018yCat} astrometry, { and the CRIRES relative radial velocity to constrain the inclination.} | {\label{conc}} We presented an improved reduction of the FORS2 astrometric measurements of the LUH 16 binary and showed that they are consistent with the HST and GeMS positions in the literature. In our previous reduction \citep{2015MNRAS}, the positions of LUH 16 were linked to the ICRF using USNO-B, and as noted by \citet{Bedin}, are inconsistent with the HST measurements at a level of 10--20~mas. This was due to a flawed transformation into the ICRF that used stars in both chips of FORS2. Before Gaia DR1, this effect could not be diagnosed and lead to a significant bias in RA, Dec of about $\pm$50~mas \citep{gaia}. It had only a negligible effect on the relative positions of the LUH 16 components via the pixel scale and rotation of the coordinate axes, however. The main reason for the inconsistency noted by \citet{Bedin} was the incorrect transformation into the ICRF with Eq.(\ref{eq:astr_ICRF}), where as the argument of functions $F_n$ we used the positions of LUH 16 $\bar x$, $\bar y$ extrapolated to the USNO-B epoch J2000.0 instead of using the positions at the actual epochs. Gaia DR2 provides the reference frame for determining absolute proper motions and parallaxes. However, we argue that very large ground-based telescopes provide us with competitive differential astrometry for faint 16--21~mag stars. In the LUH 16 field analyzed in this paper, the relative proper motions and parallaxes were determined with FORS2 with an average precision of 0.18~mas/yr and 0.13~mas, respectively, for 16--21~mag stars in common with DR2. In spite of the short one-year duration of FORS2 observations, this precision is about twice better than the corresponding values 0.48~mas/yr and 0.31~mas in the DR2 catalog. The good astrometric performance of FORS2 is due to the higher signal-to-noise ratio in the position estimation. We refined the individual dynamical masses of LUH 16 to $33.5\pm 0.3\, M_{Jup}$ (component A) and $28.6\pm 0.3\,M_{Jup}$ (component B), which corresponds to a relative precision of $\sim$1\% and is three to four times more precise than the estimates of \citet{Garcia}. We found that{ a minor bias with $1\sigma$ of} the ESO-R resolved astrometry based on a 1984 photographic plate { leads to a change in the dynamical masses that does not exceed the random error} of our determination. The { exact characterization of any residual bias in mass} will be resolved in the future when the high-precision astrometry will cover a larger portion of the orbit. Because of the importance of LUH 16 as an extremely well-studied system of nearby brown dwarfs, the refinement of the parallax, orbital, and dynamical mass parameters will continue, with the purpose of increasing the knowledge of the physical parameters of the system. We expect that the astrometric data set presented here will contribute significantly to this process. | 18 | 8 | 1808.07835 |
1808 | 1808.03627_arXiv.txt | We study the cosmic MeV neutrino background from accretion disks formed during collapsars and the coalescence of compact-object mergers. We provide updated estimates, including detection rates, of relic neutrinos from collapsars, as well as estimates for neutrinos that are produced in mergers. Our results show that diffuse neutrinos detected at HyperK would likely include some that were emitted from binary neutron-star mergers. The collapsar rate is uncertain, but at its upper limit relic neutrinos from these sources would provide a significant contribution to the Cosmic Diffuse Neutrino Background. | Accretion disks surrounding black holes (BH) or hypermassive neutron stars (HMNS) are likely the final fate of the coalescence of a neutron star (NS) with a compact object (BH or NS) \cite{Lee1999,Rosswog:2005su,Foucart:2014nda,Fryer:2015uia,Lehner:2016lxy}. Accretion disks are also formed during rare supernovae that have significant rotation, termed collapsars \cite{MacFadyen:1998vz,OConnor:2010moj,Ott:2010gv,Sekiguchi:2010ja}. During these events, much of the gravitational energy is released as neutrinos. The neutrinos are interesting not only because of their key role in the setting of the electron fraction and subsequent synthesis of elements, e.g. \cite{Surman_Mclaughlin04,Surmanrprocess,Surman:2008qf,Caballero:2011dw,Just:2014fka}, or their suspected contribution to the triggering of long duration gamma ray bursts (see e.g. \cite{Woosley:1993wj,Popham1999,Kneller:2004jr,Murguia-Berthier:2016fys}), but also because they are one of the signals that come from these multi-messenger objects. Even from a small number of neutrinos (just like in the SN1987 case \cite{Hirata:1987hu}), much can be gleaned about the central engines of these objects. Neutrinos emitted from these accretion disks are expected to be in the energy range of MeV. It is well known that astrophysical MeV neutrinos could be registered at existing facilities such as SuperKamiokande (SK) \cite{SK} and SNOLAB \cite{Descamps:2015hva}. The prospects of detection in larger facilities like the proposed Hyperkamiokande, UNO, DUNE, JUNO and TITAND \cite{hk,uno,Acciarri:2015uup,An:2015jdp,Titand}, are even more promising \cite{Kistler:2008us}. There are two basic strategies for detecting these MeV neutrinos, either a direct detection from an object that is sufficiently close to produce a substantial flux at earth, or a detection of the cosmic MeV neutrino background. The latter is formed by the accumulation of neutrinos from such objects over time. The consideration of the cosmic MeV neutrino background (CMNB) from all types of extra galactic sources (supernovae, collapsars, binary mergers) enhances our chances of detection and opens a window to neutrino physics at cosmological scales. A detection of the CMNB will provide insights to the star formation history, initial mass function, cosmic metallicity, and event rates (see e.g. \cite{Nakazato:2015rya,Davis:2017mbq}). While many types of events can contribute to the CMNB, two types of events have been explored most extensively. Due to its promising prospects for detection, the supernova relic neutrino (SRN) background has been widely studied (e.g \cite{Kaplinghat:1999xi,Beacom:2003nk,Strigari,Ando2003,Ando2004,Lunardini2006,Lunardini:2010ab}, for reviews see e.g. \cite{Beacom:2010kk,Lunardini:2010ab}). SRN searches at SK have significantly improved upper limits, and they are now very close to predictions \cite{Malek:2002ns,Bays:2011si,Zhang:2013tua}. The next most studied contribution to the CMNB is that of relic neutrinos from failed supernovae (or unnovae). Theoretical fluxes \cite{Lunardini2009failed,Keehn:2010pn,Yang:2011xd,Priya:2017bmm,Nakazato:2015rya} have been found to be comparable to that of supernovae \cite{Nakazato:2008vj}. In this paper we add to previous CMNB studies by considering the neutrino background due to accretion disks from compact object mergers and supernovae (collapsars). We use updated models to extend previous work on collapsars, for example that of \cite{Nagataki:2002bn} which found optimistic detection prospects for TITAND, using a neutrino background determined from the collapsar model in \cite{Nagatakicounts}. We also make the first determination of the diffuse neutrino background from compact object mergers. In both scenarios, matter surrounding the remnant black hole or hypermassive neutron star is hot and will emit considerable numbers of neutrinos. The study of the accretion tori allows for a determination of neutrino emission, \textit{after} black hole formation, of collapsars and mergers. By considering two different accretion disk models, discussed later, we investigate the effect that the accretion rate and the BH spin have on the neutrino spectra, on the relic background, and on the associated number of neutrinos reaching Earth's detectors. The derivation of an accretion disk diffuse flux relies on two components: the neutrino spectra emitted in one of the above scenarios and the cosmological rate at which these events occur. In both collapsars and mergers, the neutrino emission can be comparable to or even larger than supernovae. In the collapsar case, simulations have shown that the neutrino emission may be larger than that of a proto-neutron star \cite{Fischer:2008rh,Sumiyoshi:2008zw}. In the case where the disk is formed after a merger, the neutrino emission from one event, although shorter in duration, can be one or two orders of magnitude more luminous than that of a supernova \cite{Cab2009}. Similar to \cite{Nagataki:2002bn} we employ steady state models of accretion disks where the disk is considered to be the result of a collapsar, but in addition we consider a dynamical model. Our estimates come from updated models which include neutrino cooling, a range of black hole spin and estimates of gravitational bending and redshifting on the part of the neutrinos. We assume that the BH has been already formed and matter, in a torus distribution, is accreted into it at a given rate. We also comment on the case of an accretion disk surrounding a hypermassive neutron star. The other component, the cosmological failed supernova and merger rates, has been revisited in the recent years, motivating also this study. From one side the detection of gravitational waves from mergers at observatories such as Advanced-LIGO, has triggered an impressive effort to estimate the merger coalescence rates \cite{OShaughnessy:2008yul,dominik}; with estimates in the range of ($10^{-6}-10^{-3}$/year per Milky Way Equivalent Galaxy (MWEG) for NS-NS mergers, and $10^{-8}-3\times10^{-5}$/year per MWEG for BH-NS mergers \cite{Abadie:2010cf}). The recent detection of a neutron star mergers, suggests a rate of $1540^{+3200}_{-1220}$ Gpc$^{-3}$yr$^-1$ \cite{TheLIGOScientific:2017qsa}. On the other side, recent \textit{Swift} gamma rays burst data \cite{Gehrels:2004aa,Butler:2007hw} has been used to provide new estimates for star formation rates \cite{Yuksel2008} and failed supernovae \cite{Yuksel:2012zy}. In this manuscript we convolve the accretion disk neutrino spectra from two different models, with current failed supernova and merger rates. In doing so, we provide an updated baseline for future studies on relic neutrinos from collapsars, the first estimates from mergers, and comparison between the two scenarios. We focus on the electron antineutrino relic flux, its contribution to the MeV neutrino background, and its possible detection at water Cherenkov facilities. Although important, we do not consider neutrino oscillations in this work. Oscillations will change the large energy contribution of the neutrino spectra resulting in a larger number flux of relic electron antineutrinos (see e.g. \cite{Ando2003,Lunardini:2012ne}). Oscillations are expected to play a significant role in mergers and collapsars, e.g. \cite{Malkus:2012ts,Malkus:2014iqa,Zhu:2016mwa} and we will discuss the role of oscillations in the accretion disk relic neutrinos in future work. This paper is organized as follows: in section \ref{disk model} we discuss the accretion torus models used and in section \ref{spectra} we present the corresponding results for the neutrino spectra. We continue by introducing the compact object mergers and failed supernovae rates used in this work and show our results for the relic neutrino flux for each scenario in section \ref{diffuse flux}. In section \ref{detection rates} we provide neutrino event rates at SK and finally in section \ref{conclusions} we conclude. | \label{conclusions} We have studied the spectra, diffuse fluxes, and detection rates in SuperK and HyperK, of neutrinos emitted by past to present black hole accretion disks. Those are considered to be one of the possible final fates of rotating core collapse supernovae, as well as of mergers of neutron stars with black holes or with other neutron stars. The models used for our study include important aspects such neutrino cooling and relativistic effects. Neutrino disk spectra depend on the mass accretion rate and the BH spin. The evolution of accretion disks is such that there is a funnel formed around the black hole. When neutrinos are emitted from that low density region they have large temperatures. The effect of these high temperatures and the torus like geometry is reflected in a hotter neutrino spectrum compared to that from spherically symmetric SN simulations (the latter used to study the failed SN spectrum). The number of failed supernovae that evolve into a disk (in a collapsar model) depends on still to be determined physics such as the nuclear matter equation of state and the progenitor initial conditions, leaving us with open questions on the mechanism of BH formation. Future simulations would shed light into the BH mass, BH spin, and accretion rate ranges that would be possible if such tori formed from unsuccessful SN. This uncertainty notwithstanding, our spectra results motivated us to study the potential contribution of these neutrinos to the relic neutrino background in the MeV range. We find that in the collapsar model, assuming an upper limit event rate that is the same as the unnova rate, this diffuse flux is comparable (larger for high mass accretion rates) to the SN one. We find that the number of neutrinos registered in SuperK (taking an energy threshold of 5 MeV) in a 5 year period from collapsars would be between 3 to 25. As discussed elsewhere (see e. g. \cite{Priya:2017bmm}), the atmospheric and reactor neutrino fluxes limit the detection energy window from $\sim$ 11 to 30 MeV in the current SuperK setup. In that range we predict that in the most optimistic collapsar model we will find about 3 counts per year. We also studied the diffuse flux and possible detection of neutrinos coming from ADs in the scenario where the torus was the result of a neutron star-compact object merger. The cosmological merger rates lead to diffuse fluxes that are at least two orders of magnitude smaller than those of SN and collapsars. However, the upgrade from SuperK to HyperK will allow for a detection of at least one of those neutrinos (in the most optimistic merger scenario) in a period of 1.75 years. It is important also to keep in mind that these results are based on merger rates for field stellar populations \cite{dominik}, but rates should be larger in globular clusters. The rates are also sensitive to parameters in the binary model and initial distributions of the binary \cite{deMink:2015yea}. A recent compilation of different predictions of NS-NS and BH-NS merger rates can be found in \cite{Abbott:2016ymx} showing that event rates may be higher than assumed here. The prospects of overcoming the current detection limitations on the detection of the CMNB are promising. Extracting relic neutrino signals in SK, with more data collected, improved efficiency, and lower threshold will be a reality in few years \cite{Zhang:2013tua}. The possibility of a megaton water-Cherenkov detector like HyperK opens the door to significant numbers of diffuse neutrinos being observed. In analyzing such a future detection, we should bear in mind that in addition to standard core collapse and failed supernovae, a few of these neutrinos may come from accretion disk supernovae and compact object mergers. | 18 | 8 | 1808.03627 |
1808 | 1808.04377_arXiv.txt | Super-massive black holes (SMBHs) with $M_{\bullet} \sim 10^9 $ M$_{\odot}$ at $z>6$ likely originate from massive seed black holes (BHs). We investigate the consequences of seeding SMBHs with direct collapse BHs (DCBHs) ($M_{\bullet}=10^{4-6}\, \mathrm{M}_\odot$) on proto-galactic disc growth. We show that even in the absence of direct feedback effects, the growth of seed BHs reduces the development of gravitational instabilities in host galaxy discs, suppressing star formation and confining stars to a narrow ring in the disc and leading to galaxies at $z \sim 6$ which lie above the local BH-stellar mass relation. The relative magnitude of cosmic and BH accretion rates governs the evolution of the BH-stellar mass relation. For typical DCBH formation epochs, $z_{\rm{i}} \sim 10$, we find star formation is inhibited in haloes growing at the average rate predicted by $\Lambda$CDM which host BHs capable of reaching $M_{\bullet}\sim 10^9 \, \mathrm{M}_{\odot}$ by $z\gtrsim6$. Slower growing BHs cause a delay in the onset of star formation; a $M_{\bullet} \sim 10^6 $ M$_{\odot}$ seed growing at $0.25$ times the Eddington limit will delay star formation by $\sim100$ Myr. This delay is reduced by a factor of $\sim10$ if the halo growth rate is increased by $\sim 0.6\, \sigma$. Our results suggest that SMBHs seeded by DCBHs and their host galaxies form in separate progenitor haloes. In the absence of subsequent mergers, higher than average cosmic accretion or earlier seed formation ($z_{\rm i} \sim 20$) are required to place the evolving BH on the local BH-stellar mass relation by $z=6$. | \label{introduction} The relationship between super-massive black holes (SMBHs) and their host galaxies is an active area of research \citep[see, e.g.][]{Kormendy2013, Hickox2014, Delvecchio2015, Bongiorno2016, Yang2018}. Several empirical correlations between the mass of a black hole (BH) and the physical properties of its host galaxy have been reported \citep[e.g.][]{Magorrian1998, Merloni2003, Merritt2006, Kormendy2009}. Of these correlations the BH mass-stellar velocity dispersion (the $\mathrm{M_{\bullet}} - \sigma$ relation) \citep[e.g][]{Ferrarese2000, Gebhardt2000, Mcconnell2011} historically gave the first clues on a feedback driven co-evolution of BHs and their host galaxies \citep{Silk1998, King2003}. Theoretical and observational studies suggest that major mergers play a fundamental role in establishing feedback and feeding cycles \citep[e.g.][]{DiMatteo2005, Hopkins2006, Kormendy2011}. Further support for the importance of mergers comes from the increased scatter in the BH-host correlations at larger redshifts which is a natural consequence of the central-limit theorem and an increasing number of BH mergers for SMBH toward low redshifts \citep{Schawinski2006, Peng2006, Hirschmann2010}. However, mergers are not the only physical processes involved; offsets in the $\mathrm{M_{\bullet}}-\mathrm{M_{Bulge}}$ relation corresponding to disc galaxies can be explained through the co-evolution of SMBHs with their disc-galaxy hosts through secular processes \citep{Volonteri2016, Simmons2017, Martin2018}. SMBHs with masses of $M_{\bullet}\sim 10^{9}\, \mathrm{M_{\odot}}$ have been observed in galaxies at high redshifts ($z\sim$ 6 -- 7) \citep{Fan2006a, Mortlock2011, Banados2018} corresponding to less than a gigayear after The Big Bang. This population of BHs must have had a rapid formation process to reach the masses observed at this early epoch. Indeed, if the growth rate of SMBHs is limited by the Eddington accretion rate \citep[see however,][]{Natarajan2012,Pacucci2017}, they must be seeded by some massive progenitor at an early epoch $z \geq 10$, prior to the onset of reionization and the shutdown of Population-III stars \citep{Paardekooper2015,Johnson2013a}. Given the e-folding nature of an Eddington limited BH-growth rate on a Salpeter time scale of $t_{\mathrm{Sal}} = 0.45 \, \eta/(1 - \eta)\; \mathrm{Gyr}$ (where $\eta \sim 0.1$ is the radiative efficiency \citep[see, e.g.][]{King2008}), varying the initial seed mass by factors of ten can have a strong impact on relaxing the constraints on the formation time in the early Universe. Consequently, various seed formation processes are discussed \citep[see, e.g. for a review][]{Volonteri2010}, including population III stellar remnants \citep{Madau2001} and the collapse of dense stellar clusters \citep{Clark2008,Yajima2016} or the direct collapse of gas through the state of a super-massive star \citep{Bromm2003,Lodato2006,Begelman2006,Begelman2010,Agarwal2012b}. The latter channel has received heightened attention due to the massive seeds it produces and the ability to grow to super-massive scales with less stringent constraints on the average accretion rate \citep{Agarwal2012b, Latif2013, Latif2014, Pacucci2015, Agarwal2016a}. \footnote{\citet{Johnson2011} simulated the radiative feedback from such a seed BH showing that the average accretion rate is very low, indicating that the feedback might off-set the advantage you gain of having a higher initial mass.} Direct collapse BHs (DCBHs) form during the collapse of pristine gas in haloes with virial temperatures of $T_{\rm v}\gtrsim 10^{4}$ K \citep{Bromm2003}. Provided a halo remains pristine and the local intensity of the Lyman-Werner radiation field is greater than the critical intensity required to dissociate any H$_{2}$ gas \citep{Agarwal2016}, cooling within the halo will only take place via atomic hydrogen. The gas temperature in such a halo will be kept at $T_{\rm g}\sim 10^{4}$ K during collapse with a Jeans Mass of $M_{\rm J}\sim 10^{6} \, \mathrm{M_{\odot}}$; preventing the fragmentation into gas clumps and stars, and leading to the isothermal collapse of a massive gas cloud into a single BH \citep{Bromm2003}, possibly via an intermediate stage of a super-massive star \citep{Begelman2010}. This process results in the formation of massive seed BHs with masses of $M_{\bullet}\sim 10^{4}-10^{6}\, \mathrm{M_{\odot}}$ at $z\sim 10 - 20$, prior to the formation of the host galaxy in the halo \citep{Agarwal2012b}. If SMBHs are truly seeded by DCBHs it would affect the early stages of galaxy evolution. Gas build up around the gravitational potential well of a DCBH through cosmological accretion and halo merging, would not only lead to further BH growth but also potentially to the gradual growth of a proto-galaxy around the BH. Besides feedback from the BH affecting the proto-galaxy, initially such a proto-galaxy would be gravitationally dominated by the mass of the BH as well. However, it is not clear how this would affect the processes of galaxy evolution, such as star formation, and the cycle of baryons in the galaxy. Recently, a first potential candidate for an observed DCBH system has been proposed (\citealt{Sobral2015}; but see \citealt{Bowler2017}). The system, called CR7, is a very bright Ly $\alpha$ emitter at $z=6.6$ with $L_{\mathrm{Ly}\alpha} \sim 10^{44} \mathrm{erg \,s^{-1}}$ \citep{Matthee2015, Sobral2015, Bowler2017}. \citet{Sobral2015} have identified CR7 as a combination of three components: Two clumps which appear to be evolved galaxies in close proximity to a third clump, which provides the vast majority of the Ly $\alpha$ flux. This third clump has been successfully modelled by \citet{Agarwal2016a} as a $M_{\bullet} \sim 4.4 \times 10^6 \, \mathrm{M}_{\odot}$ BH formed through direct collapse around $z\sim20$. Recent work has shown either an active galactic nucleus (AGN) or a young starburst population are also likely explanations for the observed characteristics of CR7 \citep{Bowler2017}. With the former being potentially seeded via the stage of a DCBH, and the latter not requiring a DCBH at all. The evolution during the initial stages of a potential DCBH-systems such as CR7 is unknown and yet likely consists of a constant interplay between star formation and BH growth. The formation of stars in proto-galaxies is driven by the accretion of gas and subsequent gravitational collapse. The star formation law relating the star formation rate (SFR) surface density in a disc galaxy to its gas surface density \citep{Schmidt1959}, once confirmed by observations of local galaxies \citep{Kennicutt1998}, has more recently been shown to extend to $z\sim1.5$ \citep[see, e.g.][]{Carilli2013}. At higher redshifts, the higher densities imply shorter cooling times (as $t_{\rm cool} \propto \rho^{-1}$). Rapid cooling means that the SFR is only limited by the total gas accretion rate and the growth rate of gravitational instabilities in a galaxy \citep{Dekel2009,Dekel2013}. Previous studies on galaxy evolution and star formation have related the global SFR to disc properties via the growth rate of instabilities \citep{Fall1980, Lacey1983, Wang1994, Schaye2004, Elmegreen2010}. For example, \citet{Elmegreen2010} modeled the growth in gas mass and turbulence driven by gas accretion onto a galaxy and found the SFR was mainly a function of the gas accretion rate. However, star formation is a local process. Star formation can only take place where the gas is unstable to gravitational collapse \citep{Wang1994}. In this context the BH may also play an important role in the stability of the disc \citep{Lodato2012}. The gas properties will change throughout the galaxy with some regions being more unstable than others. Indeed, \citet{Schaye2004} found that if disc galaxies are rotational supported against collapse this will be particularly at large radii, limiting the radial extent of star formation to within some truncation radius. A further complicating factor for the growth of proto-galaxies around DCBHs is that the hosting halo is generally in the vicinity of a more massive halo it is likely to merge with at a later stage during its evolution \citep{Agarwal2014a}. During the satellite-phase the provision of fuel for star formation will cease due to stripping processes in the environment \citep{VanDenBosch2008}. The growth of the host galaxy will thus be affected and in turn the path to the locally observed BH-galaxy correlations. The aim of this paper is two-fold, we want to model the stabilising effect of DCBHs on the gaseous disc in proto-galaxies and their impact on the onset of star formation, and based on these models present arguments for the evolution of DCBHs toward locally observed correlations with host galaxies. First, we lay out the star formation model we use which relates star formation rate to disc instabilities (section~\ref{starform}). The model is first introduced by discussing a non-evolving case in section~\ref{Non-evolving} before being fully explored in section~\ref{Evolving} with evolving the halo and stellar mass. Finally, we discuss the implications of the model for massive seed hosting galaxies and the onset of star formation within them (section~\ref{conclusions}). Throughout the paper a $\Lambda$CDM Universe is assumed with $\rm H_{0} = 70 km\, s^{-1}\, Mpc^{-1}$, $\Omega_{\mathrm{m, \, 0}} = 0.27$ and $\Omega_{\mathrm{\Lambda, \, 0}} = 0.73$. | \label{conclusions} We use an analytical model to investigate the effect of a DCBH seed on the stability of proto-galaxy discs and the resulting suppression of star formation. We look at how the Toomre and tidal stability parameters profiles of an exponential disc change due to the presence of a BH in the centre of the system and link the stability of the disc to the star formation rate. We show how the BH has a gravitationally stabilising effect on the inner region of the disc which increases the star formation timescale locally and limits the region of the disc where star formation can occur, decreasing the modelled SFR. We also model the growth of a galaxy around a seed BH to investigate how the interplay of cosmological accretion, accretion onto the BH and the stabilizing effect of the BH can be important in determining the circumstances under which stars can form. After the initial onset of star formation, we find that the radial extent of the star forming region remains relatively constant. Under the assumption of stars staying on circular orbits and not migrating in the disc, the process of forming stars increases the local surface density ($\Sigma_{\rm g} + \Sigma_{\star}$). This increases the self-gravity of the disc locally and decreases the effect of tidal forces on the gas. Removing the support from the tidal shear against gravitational collapse then leads to the further formation of stars in this same region. Following a short period beginning at the onset of star formation (while the stellar mass is still negligible), all subsequent star formation in the disc is largely confined to the region where stars have already formed. As stability increases in the presence of a massive BH, the radial extent of the region where stars can form narrows and the total SFR is reduced. The radial extent of the region where stars can form in the model disc is small ($\sim 100$ pc) due to the disc properties at $z=6$, even in the absence of a BH. For the evolving model with a formation redshift at $z_{\rm i}=10$, we calculate the angular size of the stellar disc in the no BH case at $z=6$ to be $\theta< 0.02$ arcseconds and note this is less than the angular resolution of the James Webb Space Telescope (JWST), even at the shortest possible wavelengths. Resolved observations of such objects at this redshift would therefore by infeasible with current instruments. The presence of a growing BH seed can greatly affect the star formation history of its host galaxy, even preventing the formation of stars entirely. Increasing the mass of the BH or the scale radius of the disc increases the stability of the disc, while increasing the disc mass decreases the stability. In the fiducial case, the disc becomes more unstable in the star forming region as the disc mass increases with the growth of the halo, resulting in SFR increasing with time. We find the sSFR in the model increases with higher BH mass and that the sSFR we calculate is higher than the observed median value at high redshift \citep{Stark2013}, particularly at times close to the onset of star formation. Our results suggest that systems hosting DCBHs should occupy the upper envelope of the sSFR distribution for any given stellar mass. Indeed, high sSFR galaxies could potential be used for the identification of DCBH hosts. As we evolve the model to lower redshifts, the discrepancy between the model sSFR and the observations decreases. Increasing the accretion rate of the BH leads to an increase in the stability of the disc at a given time as the BH mass increases and the disc mass decreases. This can lead to a delay in the time where the disc first becomes unstable and forms stars. This delay in the onset of star formation is not only dependent on the BH growth rate and seed mass but also the growth of the disc and halo. As halo growth rates are higher at high redshift, the delay is also a function of the formation redshift of the BH. For a sufficiently high BH accretion rate and seed mass, the disc can be prevented from ever forming stars. At the lowest halo growth rates and high BH accretion rates, even models with early formation times have no stars forming. Such a low halo growth rate is typical of satellite galaxies \citep[see, e.g.][]{DeLucia2012}. This suggests the chance of a SMBH forming with no stellar disc counterpart is more likely in satellite galaxies. Indeed, this would also occur if an in-fall event were to occur prior to the onset of star formation. We find that the halo in which a seed is born at $z=10$ is prevented from having significant star formation if the BH grows at the Eddington limit. If a seed BH is to grow at the rate required to increase in mass by $\gtrsim 3$ orders of magnitude between $z\sim10$ and $z\sim6$, star formation in its host is suppressed, placing such a system above the BH-stellar-mass relation. This suggests that DCBH galaxies will move towards the local BH-stellar mass relation via potential mergers with already evolved galaxies without massive BHs and not self-regulated co-evolution. Alternatively, this discrepancy can be resolved if either the formation of the DCBH is pushed to higher redshift ($z\sim20$) or if the evolution of the BH-galaxy system takes place in haloes with higher than average growth rates. Though we do not model the feedback from the accreting BH we acknowledge that this would change the star formation and BH growth histories \citep{Schawinski2006, Latif2018}. BH feedback would heat and eject gas in the disc, acting to stabilise it, reducing the star formation rate in the model. The process of stabilising the disc through BH feedback would complement the gravitationally stabilising effect of the BH, delaying the onset of star formation further and decreasing the area of the disc able to form stars. This does not take into account the inclusion of ``positive feedback'' \citep{Gaibler2012}, where the inducing of star formation through jets leads to an increase in the SFR. However, this induced star formation would take place at large radii, meaning the inner region close to the BH would still be void of stars. | 18 | 8 | 1808.04377 |
1808 | 1808.04188_arXiv.txt | { We have derived a modified Lane-Emden equation for the Starobinsky model in Palatini gravity which is numerically solvable. Comparing the results to the ones provided by General Relativity we observe a significant difference depending on the theory parameter for the $M-R$ relations. } \PACS{ {04.50.Kd, 04.40.Dg, 97.60.Jd}{} % } % | \label{sec:introduction} The shortcomings \cite{Copeland:2006wr,Nojiri:2006ri,Capozziello:2007ec,Carroll:2004de,Sotiriou:2008ve} of General Relativity (GR) \cite{Einstein:1915ca,Einstein:1916vd} make the search for some other proposals describing the gravitational phenomena necessary and appealing. The dark matter idea \cite{Cap_Laur, Cap_Far}, inflation \cite{Starobinsky:1980te,Guth:1980zm}, the fact of the late-time cosmic acceleration \cite{Huterer:1998qv, Sami} with an explanation in the form of the exotic fluid called dark energy \cite{Copeland:2006wr,Nojiri:2006ri,Capozziello:2007ec} are just the most widespread problems which we face. Looking for a generalization of Einstein's theory is additionally supported by the fact that GR is non-renormalized while adding extra high curvature terms seems to improve the situation \cite{stelle}. This is why Extended Theories of Gravity (ETG) \cite{Cap_beyond, cap_invar} have gained a lot of attention. However, many extensions introduce ghost-like instability. Nonetheless, attacking the Hilbert-Einstein action appears in many different ways: assumption on the ``non-constancy'' of the Nature constants \cite{Dabrowski:2012eb, Leszczynska:2014xba, Salzano:2016pny}, minimally or non-minimally coupled scalar fields added to the Lagrangian \cite{brans, Bergmann}, or more complicated functionals than the simple linear one used in GR, for example $f(R)$ gravity \cite{buchdahl, Starobinsky:1980te}. The extra geometric terms coming from the latter approach could explain not only dark matter issue \cite{cap_jcap, cap_not} but also the dark energy problem. Since the field equations also differ from the Einstein's ones, they usually provide different behavior of the early Universe. One can formulate $f(R)$ gravity in different ways: in the metric approach, \cite{Sotiriou:2008rp,Carroll:2004de,Sotiriou:2008ve,DeFelice:2010aj,Will:1993te} Palatini one \cite{Palatini:1919di,Sotiriou:2008rp,Capozziello:2011et,Ferraris:1992dx} as well as hybrid \cite{hybrid}. We will focus on the Palatini approach in this work. The Palatini approach provides modified Friedmann equations \cite{alle_bor1, alle_bor2, alle_bor3} that can be compared with the observational data \cite{bor_kam, BSSW1, BSSW2, BSS, SSB, wojn_galax}. It shows the potential of the Palatini formulation and it is still applied to gravitational problems \cite{lavinia, roshan}. Moreover, there have appeared possibilities of observable effects in microscopic systems providing constraints on models parameters \cite{ronco,Lobo:2014nwa}. There are also disadvantages reported: lack of perturbative approach \cite{Flanagan:2003rb}, conflict with the Standard Model of particle physics \cite{Iglesias:2007nv, Olmo:2008ye}, the algebraic dependence of the post-Newtonian metric on the density \cite{Olmo:2005zr, Sotiriou:2005xe}, and the complications with the initial values problem in the presence of matter \cite{Ferraris:1992dx, Sotiriou:2006hs}, although that issue was already solved in \cite{olmo_sanch}. A similar discussion was performed in \cite{vignolo} where it was shown that the initial value problem is well-formulated in presence of the standard matter sources while the well-posedness of the Cauchy problem should be considered case by case: the Starobinsky one, which we are interested in, belongs to the well-posed class of models. There are also additional arguments showing that treating the extra terms in the fluid-like manner provides limitations \cite{zaeem, zaeem2, zaeem3}. However, as it was shown in \cite{olmo_tri}, higher curvature corrections do not cause the above mentioned problems. Interestingly, the effective dynamics of Loop Quantum Gravity as well as brane-world cosmological background histories may be reproduced by the Palatini gravity giving the link to one of the approaches to Quantum Gravity \cite{olmo_singh, olmo_brane}. There are also astrophysical aspects of Palatini gravity, for instance black holes were considered in \cite{diego1,diego2,diego3,diego4,diego5,diego6}, and also wormholes \cite{worm,diego7,diego8} and neutron stars \cite{kain,reij,pano,anab, anet}. Our concern is related to the last objects in Palatini gravity, especially that the recent neutron stars' merger observation \cite{gw} will provide the possible confrontation of gravitational theories. Neutron stars seem to be perfect objects for testing theories at high density regimes: there are claims that using General Relativity in the case of strong gravitational fields and in the case of large spacetime curvature \cite{controversialmag, eksi, berti} is an extrapolation. Stars in Palatini gravity were considered in \cite{Barausse:2007ys, barau, barau2, pani, sham} where it was claimed that there exist surface singularities of static spherically symmetric objects in the case of polytropic equation of state which can lead to infinite tidal forces on the star's surface. An argument against that claim was introduced in \cite{olmo_ns}: the problem is caused by the particular equation of state which should not be used at the surface. Moreover, in \cite{fatibene} it was indicated that the polytropic equation of state is nothing fundamental but rather an approximation of the matter forming a star. The another important point was mentioned that Palatini gravity should be interpreted according to the Ehlers-Pirani-Schild (EPS) approach \cite{eps, mauro, fatibene1}, which we are going to follow in this work. That means that a conformal metric is the one responsible for the free fall in comparison to metric which was used in \cite{Barausse:2007ys}. As shown in \cite{fatibene}, in this case the singularities are not generated and polytropic stars can be obtained in the Palatini framework. Using this result, we are going to study non-relativistic stars with the polytropic equation of state in $f(\hat{R})$ Palatini gravity. As a working example we will use the Starobinsky model, that is, $f(\hat{R})=\hat{R}+\beta\hat{R}^2$ \cite{Starobinsky:1980te}. In order to do it for the wide class of the stellar objects, we will write down the modified Lane-Emden equation obtained from the generalized Tolman-Oppenheimer-Volkoff (TOV) equation which we studied in the context of star's stability in \cite{anet}. That will allow to examine further different types of stars since the equations describing them can be solved numerically. There are already works considering modified TOV equations \cite{AltNS1, AltNS2, AltNS3, AltNS4, AltNS5, AltNS6, AltNS7, aw1, aw2} as well as ones providing generalization of the Lane-Emden \cite{riazi, capH, saito, andre, sak, koy}. We are using the Weinberg's \cite{weinberg} signature convention, that is, $(-,+,+,+)$, with $\kappa=-8\pi G/c^4$. | In contrast to the existing works on Palatini stars, we have used the EPS interpretation of the theory which provides different TOV equations. Together with the previous studies on neutron stars \cite{anet}, galaxy rotation curves \cite{wojn_galax}, and cosmology \cite{mauro, fatibene1, fatibene2} the current proposal has added new arguments in favor of Palatini gravity under the EPS formulation. We have derived the Lane-Emden equation coming from the Palatini modified equations describing the relativistic stellar object. Apart from the quadratic term in $\theta^n$ on the right-hand side it resembles the modified equations obtained already in the literature \cite{riazi, capH, saito, andre, sak, koy}. The numerical solutions of the equation (\ref{em}) pictured in the figures (\ref{fig.1}), (\ref{fig.2}) and (\ref{fig.3}) definitely shows that we deal with larger stars together with increasing the parameter $\alpha$ for $n=1$. The case $n=3$ differs a lot: for the positive parameter $\alpha$ the situation is similar like for $n=1$ while for negative values (see the curves in the bottom of the figure (\ref{fig.3})) is opposite in the case of the radius: the star is larger with respect to decreasing $\alpha$ while masses tend to decrease. Moreover, independently of the type of a star, because of the conformal transformation the case $\alpha=-\frac{1}{2}$ is excluded while below $\alpha=-\frac{1}{2}$ there are unphysical profiles. The masses are significantly larger than in GR case, especially for bigger values of the parameter in both cases. Decreasing $\alpha$ one obtains smaller masses where in the case of negative values of the parameter we deal with masses smaller than the ones predicted from GR. We have not considered the masses for the case $n=1.5$ because the numerical solutions suffered by the extrapolation procedure and we do not treat the results reliable. We should also remember that in all cases we have used the simplified mass formula (\ref{simMass}). Although our studies should be viewed as a toy model, we consider it as a first step to the more accurate stellar description which we leave for the future projects. In this sense, the recent finding of a mapping between Palatini theories of gravity and GR \cite{Afonso:2018bpv,Afonso:2018mxn} may be helpful for this analysis. Work along these lines is currently underway. | 18 | 8 | 1808.04188 |
1808 | 1808.08968_arXiv.txt | { Many motivated extensions of the Standard Model predict the existence of cosmic strings. Gravitational waves originating from the dynamics of the resulting cosmic string network have the ability to probe many otherwise inaccessible properties of the early universe. In this study we show how the spectrum of gravitational waves from a cosmic string network can be used to test the equation of state of the early universe prior to Big Bang Nucleosynthesis~(BBN). We also demonstrate that current and planned gravitational wave detectors such as LIGO, LISA, DECIGO/BBO, and ET/CE have the potential to detect signals of a non-standard pre-BBN equation of state and evolution of the early universe (e.g., early non-standard matter domination or kination domination) or new degrees of freedom active in the early universe beyond the sensitivity of terrestrial collider experiments and cosmic microwave background measurements. } | \label{sec:intro} Remarkable progress has been made in understanding the universe through detailed observations of the electromagnetic radiation emitted by the cosmos. These measurements, spanning a range of frequencies from radio to gamma-ray~\cite{Hill:2018trh}, have led to the $\Lambda$CDM model of cosmology in which the universe is currently dominated by dark energy and cold dark matter with smaller components of baryonic matter and radiation~\cite{Aghanim:2018eyx}. Extrapolating the $\Lambda$CDM model back in time suggests that the very early universe was dominated by radiation in the form of photons and other relativistic particles. This extrapolation is strongly supported by measurements of the cosmic microwave background~(CMB), corresponding to the photons that escaped after recombination when the radiation temperature was about $0.3\;\text{eV}$~\cite{Aghanim:2018eyx}. The success of Big Bang Nucleosynthesis~(BBN) in predicting primordial light element abundances gives additional convincing evidence for early radiation domination~(RD) up to temperatures close to $T \simeq 5\,\mev$~\cite{Kawasaki:1999na,Kawasaki:2000en,Hannestad:2004px}. Going even further back, the observed flatness and uniformity of the cosmos and the spectrum of density perturbations suggest that this radiation era was preceded by a period of rapid expansion such as inflation~\cite{Guth:1980zm,Linde:1981mu,Albrecht:1982wi,Akrami:2018odb}. Very little is known empirically about the state of the universe between the end of inflation and the start of BBN~\cite{Boyle:2007zx}. A minimal assumption is that inflation was followed by reheating to a very hot radiation phase with temperature $T \gg {\rm TeV}$, which then cooled adiabatically until giving way to the recent matter and dark energy phases. We refer to this scenario, with only radiation domination~(RD) over many orders of magnitude in temperature between reheating and the matter epoch, as the standard cosmology~\cite{Kolb:1990vq}. While this assumption is made frequently (and often implicitly), it has not been tested directly. Furthermore, non-standard cosmological scenarios with an extended period of domination by something other than radiation between inflation and BBN have strong motivation from many perspectives, including dark matter, axions, string compactification, reheating, and baryogenesis~\cite{Moroi:1999zb, Salati:2002md, Boyle:2007zx, Chung:2007vz, Gelmini:2008sh, Visinelli:2009kt,Erickcek:2015jza, Dutta:2016htz,Giblin:2017wlo, Visinelli:2017qga, Visinelli:2017imh, Visinelli:2018wza, Poulin:2018dzj, Dutta:2018zkg, Redmond:2018xty, Nelson:2018via}. Testing the paradigm of pre-BBN cosmology is therefore of great significance in advancing our understanding of the universe. Gravitational waves~(GWs) may provide a means of looking back in time beyond the BBN epoch and probing the universe in its very early stages~\cite{Allen:1996vm, Boyle:2007zx,Cui:2017ufi,Caprini:2018mtu}. The observation of binary mergers by the LIGO/Virgo collaboration has already given further support to the $\Lambda$CDM cosmology~\cite{Abbott:2017xzu}, although -- and this is important to the motivation of this work -- the GWs observed were created only relatively recently. Opportunity to look even further back in time with GWs arises because, in contrast to photons, GWs free-stream throughout the entire history of the cosmos. Indeed, GWs emitted as far back as inflation could potentially be detected by LIGO/Virgo~\cite{Aasi:2014mqd} or proposed future detectors such as LISA~\cite{Audley:2017drz}, BBO/DECIGO~\cite{Yagi:2011wg}, the Einstein Telescope~(ET)~\cite{Punturo:2010zz,Hild:2010id}, and Cosmic Explorer~(CE)~\cite{Evans:2016mbw}. A stable and predictable source of primordial GWs is needed if they are to be used to test very early cosmology. Two promising and well-motivated potential sources are cosmic strings~\cite{Vilenkin:1984ib,Caldwell:1991jj,Hindmarsh:1994re,Vilenkin:2000jqa} and primordial inflation~\cite{Grishchuk:1974ny,Starobinsky:1979ty,Allen:1987bk}. The application of inflationary GWs to probe the expansion history of the universe was studied in Refs.~\cite{Turner:1993vb,Seto:2003kc,Nakayama:2008wy,Nakayama:2008ip} and for non-standard histories in Refs.~\cite{Giovannini:1998bp,Riazuelo:2000fc,Sahni:2001qp,Tashiro:2003qp}. However, current limits from CMB isotropy typically push the inflationary stochastic GW spectra below the sensitivity of current and next generation detectors~\cite{Guzzetti:2016mkm,Ananda:2006af,Smith:2005mm} (although see Ref.~\cite{Lasky:2015lej}). Thus, we focus on GWs from cosmic strings in this work. Cosmic strings are approximately one-dimensional objects of macroscopic length that arise nearly generically in theories of physics beyond the Standard Model~(SM). Specific examples include topologically-stable field configurations in theories with a spontaneously broken $U(1)$ symmetry~\cite{Nielsen:1973cs,Kibble:1976sj}, as well as fundamental or composite strings in superstring theory~\cite{Copeland:2003bj,Dvali:2003zj, Polchinski:2004ia,Jackson:2004zg,Tye:2005fn}. Their macroscopic properties are mostly characterized by their tension (energy per unit length) $\mu$, which is typically on the order of the square of the energy scale of new physics that gives rise to them, and directly constrained by CMB measurements to $G\mu < 1.1\times 10^{-7}$~\cite{Charnock:2016nzm}, where $G$ is Newton's constant. Cosmic strings emit gravitational radiation as part of their cosmological evolution~\cite{Vilenkin:1981bx,Vachaspati:1984gt,Turok:1984cn,Burden:1985md}. After formation, cosmic strings are expected to quickly reach a scaling regime in which their net energy density tracks the total cosmological energy density with a relative fraction $G\mu$~\cite{Albrecht:1984xv,Bennett:1987vf,Allen:1990tv}. This regime consists of a small number of Hubble-length long strings and a collection of many closed string loops. As the universe expands, the long strings intersect and intercommute to form new loops, while the existing loops oscillate and emit radiation, including GWs. This continual transfer of energy from long strings to loops to radiation is essential for the string network density to track the total energy density of the universe, rather than becoming dominant like other topological defects such as monopoles~\cite{Zeldovich:1978wj,Preskill:1979zi} and domain walls~\cite{Zeldovich:1974uw}. In particular, the presence of cosmic strings with small $G\mu \ll 1$ need not disrupt the standard cosmology. For many classes of cosmic strings, the dominant radiation emission is in the form of GWs. This is true for ideal Nambu-Goto strings, many types of cosmic strings emerging from superstring theory, and possibly those created by local $U(1)$ symmetry breaking~\cite{Olum:1999sg,Moore:2001px} (although see Refs.~\cite{Vincent:1997cx,Bevis:2006mj,Figueroa:2012kw,Helfer:2018qgv} for arguments that local strings emit mainly massive vector and Higgs quanta instead). In contrast, cosmic strings derived from global symmetry breaking are expected to radiate mainly to light Goldstone quanta~\cite{Srednicki:1986xg,Vilenkin:1986ku,Damour:1996pv,Cui:2008bd,Long:2014mxa}, with much weaker emission to GWs. We focus on cosmic strings that radiate significantly to GWs through loop formation and emission in this work. The GW frequency spectrum from a cosmic string network is sensitive to the evolution of the cosmos when the GWs were emitted. In any given frequency band observed today, the dominant contribution to the signal comes from loops emitting GWs at a specific time in the early universe~\cite{Caldwell:1991jj,Allen:1996vm,Cui:2017ufi}. As a result of this frequency-time relation, we show that the cosmological equation of state leaves a distinct imprint on the frequency spectrum of GWs from cosmic strings. Moreover, the portion of the spectrum from loops formed and emitting during RD has a distinctive nearly flat plateau with a substantial amplitude over many decades in frequency. Measuring the GW signal from a cosmic string network over a range of frequencies could therefore provide a unique picture of the very early universe that could potentially expand back before the era of BBN. The outline of this paper is as follows. After this introduction, we review cosmic string scaling and derive the GW frequency spectrum from a string network in Section~\ref{sec:spectrum}. We also exhibit the relationship between the GW spectrum and the loop emission rates and formation times, and apply these to the concurrent background cosmology. In Section~\ref{sec:cosmomap} we show how this relationship together with the anticipated sensitivities of current and planned GW detectors can be used to test the standard cosmological scenario as well as deviations from it, including large numbers of additional (massive) degrees of freedom and modified equations of state. Some of the challenges to detecting these GW signals and identifying them as coming from cosmic strings, and ways to overcome them, are discussed in Sec.~\ref{sec:detect}. Finally, Section~\ref{sec:conc} is reserved for our conclusions. The results in this paper expand upon those of our previous study in Ref.~\cite{Cui:2017ufi}. Relative to the work, we present in great detail the time-frequency connection of cosmic string GWs and its relation to the background cosmology. We also expand significantly on the experimental sensitivity of GW probes to new degrees of freedom active during early universe with presence of cosmic string dynamics, and extend our study of standard and modified cosmological histories. | \label{sec:conc} Standard cosmology maintains that an era of radiation domination began in the early universe and was followed by matter domination, which then ultimately yields to an increasing acceleration era dominated by the cosmological constant. This framework is well tested and is found to be self-consistent by a multitude of experimental probes including measurements of the CMB, supernovae, large-scale structure, and abundances of nuclei as predicted by BBN epoch. Unfortunately, the traditional experimental probes reach back only as far as BBN, which corresponds to temperatures below only about 5 MeV. There are many ideas for new physics above 5 MeV that disrupt the standard cosmology, whether it be through a different scaling phase other than radiation domination (e.g., matter or kination domination), or through extra degrees of freedom beyond the known Standard Model ones that substantially modify radiation era dynamics. Therefore, testing for new physics, and an altered cosmological evolution at temperatures greater than $5\mev$, requires new methods. The potential answer is gravitational waves, whose very early origins pass safely through recombination and BBN, which scrambles the otherwise powerful CMB probes and BBN constraints. A strong early universe source of GWs must be present in order to probe the effects that cosmological evolution can have on it. Furthermore, this source must have a reasonably well understood emission spectrum -- analogous to the \textit{standard candles} of supernovae -- with which to propagate through various assumed cosmological histories and compare with observational data. A prime candidate for this is cosmic strings, whose network formation and emission spectrum has been well studied and understood, particularly featuring a long flat plateau at high frequency during standard radiation dominated era. Another reason cosmic stings are useful GW sources to consider is that they are generically expected in a wide variety of high-scale theories of particle physics, ranging from unified field theories containing abelian factors to fundamental string theory. We have assumed the existence of cosmic strings in the early universe and have worked out the GW relic abundance vs.\ frequency spectrum for many different cosmic string tensions $G\mu$. We have reiterated previous results in the literature that GWs are an excellent way to constrain and find evidence for cosmic strings even within standard cosmological evolution (see Fig.~\ref{fig:GWGmuplot}). In addition, and what is central to our study, the GWs from cosmic strings enable the probing of modifications of early universe cosmology in regimes that no other probe can. We studied two main ways that early universe cosmology can change. First, we studied the effect of having a very large number of additional degrees of freedom present in the spectrum at high energy. If the degrees of freedom are present down to temperature $T_\Delta$ one finds that there is a frequency $f_\Delta$ above which the GW energy density is altered compared to the expectations of standard cosmology (with SM degrees of freedom). The signal for the onset of a high number of degrees of freedom is therefore standard $\Omega_{\rm GW}(f)$ vs.\ $f$ for cosmic strings up to $f_\Delta$ and then a fall-off for $f>f_\Delta$ compared to expectations. Fig.~\ref{fig:dofexample} shows the effect in the $\Omega_{\rm GW}(f)$ vs.\ $f$ plane. A second example is GWs from cosmic strings evolving in a non-standard phase, either of an early matter domination phase $(n=3)$ or an early kination $(n=6)$ phase. The early matter phase may be due the presence of a large density of heavy new physics states that later decay bringing the universe back to radiation era, which is needed to satisfy BBN constraints. In other words, the universe transitions from radiation domination at very high temperatures to matter domination (at $T$ comparable to mass of long-lived heavy new particles) and then back to radiation domination (by decay of said particles) before the onset of BBN. The kination ($n=6$) phase arises from oscillating scalar moduli in the early universe, which then decay. This leads to a cosmological history of very early radiation domination to kination domination (oscillation energy dominating) and back to radiation (by decay of the moduli). The ability to probe these alternative cosmological histories well by cosmic strings partly derives from the property that cosmic strings rapidly enter a scaling regime, which means their energy density scales with scale factor $a$ exactly the same as the dominant energy density of the universe. If there is an early matter domination phase then GW energy density scales like $a^3$ during that phase, and if there is an early kination phase, cosmic strings will scale like $a^6$ during that phase. The scaling behavior of cosmic strings means that the energy density of the GWs emitted will be altered substantially through its non-standard redshifting. Our numerical work shows the effect quantitatively, which leads to a sharp fall-off in $\Omega_{\rm GW}(f)$ at high frequency $f$ (corresponding to the new phase era) if there is early matter domination, and a sharp rise in $\Omega_{\rm GW}(f)$ if there is an early kination phase. The results are illustrated in Fig.~\ref{fig:gwmod}. GW detectors have given us a window to early universe cosmology complementary to any other probes previously developed. We have argued that a strong and well-understood source of GWs in the early universe could give us unprecedented ability to probe cosmological energy evolution of the early universe far earlier than previously attainable. We have also demonstrated that cosmic strings, if they exist, would be excellent standard candles to achieve these aims. | 18 | 8 | 1808.08968 |
1808 | 1808.04822_arXiv.txt | Inspiraling and merging binary neutron stars (BNSs) are important sources of both gravitational waves and coincident electromagnetic counterparts. If the BNS total mass is larger than a threshold value, a black hole ensues promptly after merger. Through a statistical study in conjunction with recent LIGO/Virgo constraints on the nuclear equation of state, we estimate that up to $\sim 25\%$ of BNS mergers may result in prompt collapse. Moreover, we find that most models of the BNS mass function we study here predict that the majority of prompt-collapse BNS mergers have $q\gtrsim 0.8$. Prompt-collapse BNS mergers with mass ratio $q \gtrsim 0.8$ may not be accompanied by either kilonovae or short gamma-ray bursts, because they unbind a negligible amount of mass and form negligibly small accretion disks onto the remnant black hole. We call such BNS mergers ``orphan''. However, recent studies have found that ${10^{41-43}(B_p/10^{12}\rm G)^2 erg\, s^{-1}}$ electromagnetic signals can be powered by magnetospheric interactions several milliseconds prior to merger. Moreover, the energy stored in the magnetosphere of an orphan BNS merger remnant will be radiated away in ${\mathcal O}(1\ \rm ms)$. Through simulations in full general relativity of BNSs endowed with an initial dipole magnetosphere, we find that the energy in the magnetosphere following black hole formation is $E_B \sim 10^{40-42}(B_p/10^{12}\rm G)^2$ erg. Radiating $\sim 1\%$ of $E_B$ in 1 ms, as has been found in previous studies, matches the premerger magnetospheric luminosity. These magnetospheric signals are not beamed, and their duration and power agrees with those of non-repeating fast radio bursts (FRBs). These results combined with our statistical study suggest that a non-repeating, {\it precursor} FRB may be the most likely electromagnetic counterpart of prompt-collapse BNSs. Detection of a non-repeating FRB coincident with gravitational waves from a BNS merger may settle the extragalactic origin of FRBs and can place constraints on the nuclear equation of state. FRBs can also initiate triggered searches for weak signals in the LIGO/Virgo data. | The LIGO and Virgo collaborations have already reported the direct detection of gravitational waves (GWs) from the inspiral and merger of five binary black holes~\cite{FirstDirectGW,Abbott:2016nmj,Abbott:2017vtc,Abbott:2017oio,Abbott:2017gyy} and one binary neutron star (BNS)~\cite{TheLIGOScientific:2017qsa} (event GW170817), that was accompanied by multiple electromagnetic (EM) counterparts~\cite{Monitor:2017mdv,GBM:2017lvd}. The consequences for astrophysics and fundamental physics from these observations are far reaching, and it is a matter of time until the detection of such compact binaries becomes routine. Merging BNSs are not only important sources of GWs, but also sources of coincident EM counterparts. These systems had long been suspected as the progenitors of short gamma-ray bursts (sGRBs)~\cite{EiLiPiSc,NaPaPi, Pacz86,Piran:2002kw,bergeretal05,Foxetal05,hjorthetal05, bloometal06,prs15,Baiotti:2016qnr,Paschalidis:2016agf}. The detection of the GW170817-counterpart GRB170817A~\cite{Monitor:2017mdv} has provided the best evidence, yet, that some sGRBs are powered by BNSs. BNSs are also sources of kilonovae/macronovae~\cite{Lattimer1974ApJ...192L.145L,Li:1998bw}. The association of kilonova AT 2017gfo/DLT17ck with GW170817~\cite{GBM:2017lvd} has verified this expectation, too. Merging BNSs may also be progenitors for {\em fast radio bursts} (FRBs) -- a new class of radio transients lasting between a few to a couple of tens of milliseconds~\cite{Lorimer:2007qn,Thornton:2013iua}. So far 36 FRBs have been detected~\footnote{See~http://www.astronomy.swin.edu.au/pulsar/frbcat for an up-to-date-catalog of FRBs.}. The discovery of the repeating FRB ``FRB121102''~\cite{Spitler:2016dmz} points to a non-catastrophic origin as opposed to a collapse or merger, which suggests that there may be at least two different classes of FRB progenitors. Several models have been proposed to explain FRBs including magnetar giant flares, coherent radiation from magnetic braking at BNS merger, blitzars, dark-matter induced collapse of neutron stars, axion-miniclusters, newborn highly magnetized neutron stars in supernova remnants, black hole--neutron star batteries, charged black hole (BH) binaries, black hole current sheets, black hole superradiance induced by plasma~\cite{Popov:2007uv,Totani:2013lia,Falcke:2013xpa,Bramante:2014zca, Tkachev:2014dpa,Mingarelli:2015bpo,Liebling:2016orx,Marcote:2017wan, Zhang:2017ndi,Nicholl:2017slv,Conlon:2017hhi,Margalit:2018bje,Most:2018abt}. Kilonovae from BNS mergers require dynamical ejection of matter during merger and/or from an accretion disk by neutrino irradiation, see e.g.~\cite{Metzger:2016pju} for a review. It is also widely accepted that BNSs can generate sGRBs, if a jet is launched by the BH-disk engine that forms following merger. Thus, in a scenario where a negligibly small disk forms, and a negligible amount of mass escapes, one may expect no sGRB and an undetectable kilonova from the BNS event. We call such ``kilonova-free'' and ``sGRB-free'' BNS mergers ``orphan''. But, are there any scenarios where such orphan BNS mergers arise? Numerical relativity simulations have shown that when the BNS total mass ($M_{\rm tot}$) is greater than a threshold mass ($M_{\rm thres}$), a BH ensues in the first millisecond after merger. In this prompt-collapse scenario a negligible amount of matter is ejected dynamically~\cite{Hotokezaka2013} (see also~\cite{Dietrich:2016hky}) and a negligible amount of matter is available to form a disk~\cite{STU1,STU2,lset08,Hotokezaka:2011dh,Hotokezaka2013}. Negligibly small disks were also reported in~\cite{Ruiz:2017inq}, where it was demonstrated that in prompt-collapse BNS mergers a jet cannot be launched as opposed to the ``delayed'' collapse scenario~\cite{Ruiz:2016rai}. For illustration we also note that ejecta masses $\sim 0.025-0.05M_\odot$ are required to explain the kilonova associated with GW170817~\cite{Coulter:2017wya,Drout:2017ijr,Shappee:2017zly,Kasliwal:2017ngb, Tanaka:2017qxj,Arcavi:2017xiz,Pian:2017gtc,Smartt:2017fuw,Soares-Santos:2017lru, Nicholl:2017ahq,Cowperthwaite:2017dyu}, while typical ejecta from equal-mass, prompt-collapse BNS mergers are ${\mathcal O}(10^{-4}M_\odot)$~\cite{Hotokezaka2013}~\footnote{Such small ejecta masses constitute $\sim 0.01\%$ of the total rest-mass and it is not clear that numerical relativity simulations have achieved such high levels of accuracy, yet.}, and disk masses ${\mathcal O}(10^{-3}M_\odot)$~\cite{Hotokezaka:2011dh}. According to~\cite{Metzger:2011bv} ejecta masses ${\mathcal O}(10^{-3}M_\odot)$ or greater are required for detectable kilonovae at the depth and cadence of the normal LSST survey with current or planned telescopes. Therefore, prompt-collapse BNS mergers may be orphan unless they take place nearby. But, are all such mergers expected to be orphan? If the binary mass ratio $q$ (defined here to be less than unity) is smaller than 0.8, then both appreciable matter may become unbound and a sizable disk onto the remnant BH may form~\cite{Rezzolla:2010fd,Hotokezaka2013,Dietrich:2016hky}. This is because for substantially asymmetric BNSs the lighter companion is tidally disrupted before merger, in contrast to near equal-mass binaries. Thus, sufficiently asymmetric, prompt-collapse BNS mergers may power both sGRBs and kilonovae. In this work we perform a statistical study to assess the astrophysical relevance of prompt-collapse BNSs, and the likelihood of orphan BNS mergers. In particular, we compute the $M_{\rm tot}$ and $q$ distribution of BNSs using the Galactic NS mass function and population synthesis models in conjunction with GW170817 constraints on the nuclear equation of state (EOS). We estimate that up to $\sim 25\%$ of all BNSs may result in prompt collapse. We also find that most models of the BNS mass function we treat predict that the majority of prompt-collapse BNSs have $q\gtrsim 0.8$. Furthermore, the larger $M_{\rm thres}$ is, the more skewed toward $q=1$ the distribution of binaries with $M_{\rm tot} > M_{\rm thres}$ becomes. Thus, most prompt-collapse BNSs may be orphan. But, does this imply no EM counterparts from such mergers? Recent work found that interactions in compact binary magnetospheres~\cite{Palenzuela:2013hu,Palenzuela:2013kra,Paschalidis:2013jsa,Ponce:2014sza,Ponce:2014hha} (see also~\cite{Hansen:2000am,McWilliams:2011zi,2012ApJ...755...80P,Lai:2012qe} for related discussions) can power $\sim {10^{41-43}(B_p/10^{12}\rm G)^2 erg\, s^{-1}}$ EM signals several milliseconds prior to merger. Here $B_p$ is the magnetic field strength at the pole of the neutron star. Moreover, following BH formation there is a significant amount of energy stored in the magnetosphere of the remnant. Studies of magnetospheres of stars collapsing to BHs~\cite{Baumgarte:2002vu,Lehner:2011aa,Dionysopoulou2013} have shown that a fraction on the order of $\epsilon = 1\%$~\footnote{The fraction is $20\%$ for electrovacuum~\cite{Baumgarte:2002vu}} of the total energy stored in a force-free magnetosphere is radiated away on a collapse timescale $\tau_{\rm FRB}$. This timescale is $\mathcal{O}(1\ \rm ms)$ for a NS. For a magnetic dipole in flat spacetime the total magnetic energy in the magnetosphere is \begin{eqnarray}\label{EB} E_{\rm B} & \sim & \int_R^{\infty}\int_0^\pi \frac{B^2}{8\pi} \left(\frac{R}{r}\right)^6 \frac{5+3\cos(2\theta)}{8} 2\pi r^2 \sin\theta dr d\theta \nonumber \\ & \sim & \frac{1}{12}B^2R^3 \sim 10^{41} B_{12}^2 R_{10}^3\ \rm erg, \end{eqnarray} implying an outgoing EM luminosity of \begin{equation}\label{LFRB} L_{\rm FRB} \sim 10^{42} \epsilon_{0.01}B_{12}^2 R_{10}^3\tau_{\rm FRB,1}^{-1}\ \rm erg\ s^{-1}. \end{equation} Here, $B_{12}=B_p/10^{12}\ \rm G$, $R_{10}$ the stellar radius in units of 10 km, $\epsilon_{0.01}$ the efficiency $\epsilon$ normalized to 0.01, and $\tau_{\rm FRB,1}$ the emission time in units of 1 ms. This outgoing luminosity matches the premerger magnetospheric luminosity. Moreover, the power and duration of these magnetospheric signals match those of observed FRBs~\cite{Mingarelli:2015bpo}. Thus, BNSs are candidates for non-repeating, {\it precursor} FRBs. Note that when two NSs merge and collapse to a BH promptly, the total energy stored in the magnetosphere is likely about the same order of magnitude as in Eq.~\eqref{LFRB}, because there is little time available to amplify the surface magnetic field through hydromagnetic instabilities as in a delayed collapse scenario~\cite{Kiuchi:2015sga}. However, compression due to the collision can amplify the magnetic field because of magnetic flux freezing. On the other hand, a large amount of the energy will quickly fall into the remnant BH. Thus, a detailed numerical relativity study of prompt-collapse BNS mergers is necessary to assess the post-merger magnetospheric energy of BNSs resulting in prompt collapse. To confirm the expectation from Eq.~\eqref{LFRB}, we perform fully general relativistic, ideal magnetohydrodynamics simulations of prompt-collapse BNS mergers. At BH formation we compute the energy stored in the magnetosphere. Assuming a $0.8\%$ radiation efficiency~\cite{Lehner:2011aa} and a millisecond emission time, we estimate an outgoing burst with luminosity $L_{\rm EM} \sim 10^{41-43} (\epsilon/0.01)(B/10^{12}\rm) G)^2$ erg/s. Thus, our simulations provide support to the idea that prompt-collapse BNSs are promising FRB sources in addition to being GW sources. To sum, BNS mergers are promising candidates for non-repeating, precursor FRBs, and such FRBs may be the most promising EM counterpart of orphan BNS mergers. The outgoing magnetospheric burst is rather isotropic~\cite{Lehner:2011aa,Palenzuela:2013kra,Paschalidis:2013jsa}, in contrast to a sGRB which is beamed, making the detection of such FRB signatures largely independent of the binary orientation. Detection of an FRB can trigger searches in LIGO/Virgo data. The discovery of coincident GWs with an FRB may settle the extragalactic origin of FRBs. Moreover, detection of an FRB from an orphan BNS merger could provide strong evidence that the merger resulted in prompt collapse to a BH, and could place constraints on the nuclear EOS, see e.g.~\cite{Bauswein:2013jpa}. The remainder of the paper is organized as follows. In Sec.~\ref{sec:Prob} prompt-collapse BNS mergers are motivated through a study of the BNS $M_{\rm tot}$ and $q$ distribution. A description of our simulations and results are presented in Sec.~\ref{sec:sims}. Our conclusions are provided in Sec.~\ref{sec:conclusion}. Geometrized units ($G=c=1$) are adopted throughout, unless otherwise specified. | \label{sec:conclusion} In this paper, we performed a statistical study of the total mass and mass ratio distribution of BNSs using the Galactic NS mass function and population synthesis models in conjunction with recent constraints on the nuclear EOS from GW170817. We find that up to $\sim 25\%$ of all BNS mergers may result in prompt collapse. Moreover, our analysis shows that most of the considered models of the BNS mass function predict that the majority of prompt-collapse BNS mergers have $q\gtrsim 0.8$, and that the larger $M_{\rm thres}$ is, the closer to unity the $q$ distribution of prompt-collapse binaries approaches. Prompt-collapse BNSs with $q>0.8$ are likely to unbind a negligible amount of mass, and form negligibly small disks onto the remnant black holes. Thus, neither detectable kilonovae nor sGRBs may accompany the GWs from such prompt collapse BNSs. We call these kilonovae- and sGRB-free BNS mergers orphan. Our statistical study suggests that most prompt-collapse BNS mergers may be orphan. We argued that premerger magnetospheric interactions and the release of energy stored in the magnetosphere of the merger remnant can match the {\it duration and power} of FRBs. Thus, BNS mergers are promising sources of detectable, non-repeating, {\it precursor} FRBs, and FRBs may be the most promising electromagnetic counterpart of orphan BNS mergers. The outgoing magnetospheric burst in these cases is rather isotropic, making the detection of coincident FRB and GW signatures possible. We have also performed magnetohydrodynamic simulations in full general relativity of different BNS configurations that undergo {\it prompt} collapse. The stars are initially seeded with a dipolar magnetic field that extends from the NS interior into the exterior. We computed the energy stored in the magnetosphere at black hole formation, and estimated the outgoing electromagnetic luminosity produced. We find a luminosity that matches those of FRBs $L_{\rm FRB}\sim 10^{41-43}B_{12}^2 \rm erg\,s^{-1}$. We close with a few caveats: First, our statistical analysis can be refined as soon as ground based GW interferometers unveil the NS mass function in BNSs; second if one is interested in the LIGO/Virgo {\it observed} mass function, the delay-time distribution should be considered, which we do not account for here; third, some conclusions in our work are based on the size of ejecta and black hole disks found in numerical relativity simulations of prompt-collapse BNS mergers. The number of such simulations is small compared to simulations of BNS mergers resulting in delayed collapse. Therefore, more high resolution simulations in full general relativity of BNSs resulting in prompt collapse are necessary to solidify the results that such mergers unbind negligible amounts of mass and form negligibly small disks onto the remnant black hole, and to find the ``critical'' mass ratio below which appreciable mass ejection and disks occur. This critical mass ratio is also likely to be equation-of-state dependent. Fourth, whether the FRB signature will be luminous enough depends on the NS surface magnetic field. We adopted a value of $\sim 10^{12}$ G, but we note that FRB-level luminosities from magnetospheric interactions are possible even from $\sim 10^{11}$ G~\cite{Palenzuela:2013hu}. Finally, with our code we are able to obtain only crude estimates of the energy in the magnetosphere. To assess the full FRB signature in the model considered here requires a code (such as that of~\cite{Palenzuela:2013hu,Palenzuela:2013kra}) that can evolve through inspiral, merger and prompt collapse to magnetosphere release, while smoothly matching the ideal magnetohydrodynamic stellar interior to a force-free exterior. Such a simulation is currently lacking and will be the subject of future work of ours. | 18 | 8 | 1808.04822 |
1808 | 1808.05223_arXiv.txt | We investigate the value of the near-infrared imaging from upcoming surveys for constraining the ellipticities of galaxies. We select galaxies between $0.5 \le z < 3$ that are brighter than expected {\it Euclid} sensitivity limits from the GOODS-S and N fields in CANDELS. The co-added CANDELS/HST V+I and J+H images are degraded in resolution and sensitivity to simulate {\it Euclid}-quality optical and near-infrared (NIR) images. We then run GALFIT on these simulated images and find that optical and NIR provide similar performance in measuring galaxy ellipticities at redshifts $0.5 \le z < 3$. At $z>1.0$, the NIR-selected source density is higher by a factor of 1.4 and therefore the standard error in NIR-derived ellipticities is about 30\% smaller, implying a more precise ellipticity measurement. The good performance of the NIR is mainly because galaxies have an intrinsically smoother light distribution in the NIR bands than in the optical, the latter tracing the clumpy star-forming regions. In addition, the NIR bands have a higher surface brightness per pixel than the optical images, while being less affected by dust attenuation. Despite the worse spatial sampling and resolution of {\it Euclid} NIR compared to optical, the NIR approach yields equivalent or more precise galaxy ellipticity measurements. If systematics that affect shape such as dithering strategy and point spread function undersampling can be mitigated, inclusion of the NIR can improve galaxy ellipticity measurements over all redshifts. This is particularly important for upcoming weak lensing surveys, such as with {\it Euclid} and WFIRST. | Weak gravitational lensing (WL) is a slight deflection of light rays from distant galaxies when they propagate through the tidal gravitational field of intervening large scale structure. The amplitude of the WL distortion can be used to map dark matter and measure dark energy by statistically quantifying the shear distortions encoded in the observed shapes of background galaxies, namely galaxy ellipticities \cite[e.g.][]{kai95, bar01}. The ellipticities of galaxies are typically distorted only about a per cent by WL \citep{tro15}, so the WL signal in individual galaxies is challenging to detect. WL measurements thus rely on averaging over a very large sample to obtain the distortions and sufficiently unbiased estimates of galaxy shapes, which in turn require a correction for the impact of the point spread function (PSF) of the telescope. In that sense, WL observations demand high quality images because it requires a large number density of resolved galaxies and high signal-to-noise ratio (SNR), while minimizing the PSF corrections and related systematic uncertainties, with well-sampled PSFs \citep{mas13, sch18}. {\it Euclid} \citep{lau11} is a survey mission designed to understand the expansion and growth history of the Universe, and is scheduled to launch in the next decade. {\it Euclid} will image 15,000\,deg$^{2}$ of sky in one broad optical band VIS spanning 550 to 920\,nm, and three additional near-infrared (NIR) bands ($Y$, $J$, and $H$). {\it Euclid} will detect cosmic shear with VIS by measuring ellipticities of $\sim30$ resolved galaxies per arcmin$^2$ with a resolution better than 0.18$\arcsec$ (PSF FWHM) with 0.1$\arcsec$ pixels. The near-infrared bands will primarily be used to derive photometric redshifts for the weak lensing sample, in conjunction with ground-based observations at visible wavelengths. The {\it Euclid} wide survey is expected to provide WL galaxy shape measurements for 1.5 billion galaxies with space-quality resolution. To measure WL through surveys, one should measure galaxy ellipticities and its uncertainty, including systematics, very accurately. In particular, it is necessary to measure the shapes of typically faint and small, distant galaxies with high-SNR observations. In this work, we demonstrate that NIR bands result in a comparable or more precise galaxy ellipticity measurement compared to optical bands for WL studies despite their worse spatial resolution (0.3$\arcsec$ compared to 0.18$\arcsec$) and pixel sampling (0.3$\arcsec$ vs. 0.1$\arcsec$ pixel scale). There are several advantages to using NIR bands \citep{tun17}; first, NIR wavelengths sample the rest-frame optical light, which traces the older stellar population (hence the bulk of stellar mass) and is less affected by dust extinction. The VIS band covers the rest-frame UV and blue wavelengths, which predominantly traces emission from star-forming regions \citep{dic00}. In particular, the shapes of galaxies as seen in the rest-frame UV are more clumped and irregularly distributed than older stellar populations. The second advantage is that galaxies in the NIR bands have an intrinsically smoother light distribution resulting in a lower shape noise than in the optical \citep{sch18}. Third, NIR images of galaxies have a higher surface brightness with more than nine times the number of source photons per pixel, based on a calculation using images in this study; this is at least partly due to the relative importance of the bulge compared to the disk as a function of wavelength. Finally, we find that the NIR bands are sensitive to a larger number density of distant galaxies than the VIS band (see Section 2). In this paper, we study the shapes of the galaxy sample expected from {\it Euclid}-quality imaging and forecast how we can improve the shape measurement by using co-added NIR images\footnote{The simulated {\it Euclid} images in this paper do not have {\it Euclid}-specific systematics dealing with dither strategy, field distortion, PSF variations and intrapixel quantum efficiency variations which will be investigated in the future. However, it should be noted that these affect both optical and near-infrared images.}. To do that, we select galaxies from {\it HST}/CANDELS (Cosmic Assembly Near-infrared Extragalactic Legacy Survey; \cite{gro11}; \cite{koe11}) observations satisfying the {\it Euclid} sensitivity limits and simulate {\it Euclid}-resolution images. The structure of this paper as follows. The sample selection using CANDELS data is introduced in Section 2. We describe the procedure of simulating {\it Euclid}-quality optical and NIR images from {\it HST} images in Section 3. In Section 4, we explain how GALFIT (Peng et al. 2010) is used to measure the ellipticity of galaxies after accounting for the PSF, and compare the ellipticities obtained from GALFIT in simulated-{\it Euclid} and CANDELS images. Finally, we present our conclusions in Section 5. | We investigate galaxy ellipticities in simulated {\it Euclid}-quality optical (VIS) and near-infrared (NIR) images constructed from {\it HST}/CANDELS co-added V+I (VI) and J+H (JH) images. We select galaxies in CANDELS GOODS-S and -N fields (covering about 340 arcmin$^2$) with photometry in I and H bands at similar depths as the planned {\it Euclid} survey. In this study, we specifically use galaxies at a redshift range of $0.5\le z<3$. After applying a SNR $\gtrapprox 10$ cut, the total number density of galaxies at $0.5\le z<3$ is comparable in the NIR and optical; however, the NIR bands select 1.4 times higher number density of galaxies relative to the optical selection at $z>1.0$, which enable us to reduce the statistical uncertainties in the shape measurements of distant galaxies. By co-adding {\it Euclid}-quality J and H-band images, which double the number of frames and the exposure time, we can generate images at 0.15$\arcsec$ pixel scale (half of the original {\it Euclid} NIR detector pixel scale) and achieve $S/N \gtrapprox 7$ at $H<24$\,mag and $\gtrapprox11$ at $H<23.5$\,mag. Using GALFIT, we measure ellipticities of galaxies in CANDELS I and H band images with 0.06$\arcsec$ pixel scale and the simulated {\it Euclid} VI and JH images with 0.1$\arcsec$ and 0.15$\arcsec$ pixel scale, respectively. We then compare ellipticities between CANDELS and simulated {\it Euclid}-quality images while considering {\it HST}/CANDELS ellipticity as the original ellipticity of a galaxy due to its superior depth and resolution. A comparison between ellipticities derived from CANDELS and {\it Euclid}-quality VIS and NIR imaging shows that both wavelength ranges provide similar performance in measuring galaxy ellipticities at all redshifts included in this study despite the worse spatial resolution and pixel sampling of the NIR imaging. When combined with the higher source density in the NIR selection, we find that the standard error in NIR-derived ellipticities is about 30\% smaller than the optical bands at $z>1.0$, which implies a more precise ellipticity measurement than in the optical alone. Since the VIS and NIR galaxy shape measurements with {\it Euclid} have different fractional contributions of the bulge and disk, a combination of the two can improve the precision with which galaxy ellipticities are measured. The next step that is required before the NIR data can be used for WL studies is to assess how the drizzling affects both the undersampled telescope PSF, and the correlated noise \citep[see e.g.][]{rho07}. However, even though the FWHM of the {\it Euclid} NIR imaging does not quite reach {\it HST} or WFIRST resolution, the NIR data provides a major advantage for weak lensing measurements compared to optical ground-based observations that typically achieve a PSF FWHM$\sim 0.6\arcsec-0.7\arcsec$ in good seeing conditions \citep[e.g.][]{kui15, man18}; while the latter provides good sensitivity to the weak lensing signal with a median redshift of the sample of z $\sim 0.85$, about half the galaxy sample will be unresolved due to the small size of galaxies, as shown in Figure~\ref{fig:re}, implying a higher statistical uncertainty in their ellipticities. In conclusion, by using co-added J+H band {\it Euclid}--quality images, we show that the galaxy sample selected at NIR wavelengths yields a more precise ellipticity measurement, especially at high redshifts. This suggests that a careful evaluation of NIR shape systematics for future weak gravitational lensing surveys, such as with {\it Euclid} and WFIRST, should be undertaken. | 18 | 8 | 1808.05223 |
1808 | 1808.00406_arXiv.txt | HD\,219134 hosts several planets, with seven candidates reported, and the two shortest period planets are rocky (4-5 $M_{\oplus}$) and transit the star. Here we present contemporaneous multi-wavelength observations of the star HD\,219134. We observed HD\,219134 with the Narval spectropolarimeter at the Observatoire du Pic du Midi, and used Zeeman Doppler Imaging to characterise its large-scale stellar magnetic field. We found a weak poloidal magnetic field with an average unsigned strength of 2.5 G. From these data we confidently confirm the rotation period of 42 days, measure a stellar inclination of 77$\pm$8$^{\circ}$, and find evidence for differential rotation. The projected obliquity of the two transiting super-Earths is therefore between 0 and 20$^{\circ}$. We employed HST STIS observations of the Ly$\alpha$ line to derive a stellar wind mass-loss rate of half the solar value ($10^{-14} M_{\odot} {\rm yr}^{-1}$). We further collected photometric transit observations of the closest planet at near-UV wavelengths centred on the Mg{\sc ii}\,h\&k lines with {\it AstroSat}. We found no detectable absorption, setting an upper limit on the transit depth of about 3\%, which rules out the presence of a giant magnesium cloud larger than $9 R_{\rm planet}$. Finally, we estimated the high-energy flux distribution of HD\,219134 as seen by planets b and c. These results present a detailed contemporaneous characterisation of HD\,219134, and provide the ingredients necessary for accurately modelling the high-energy stellar flux, the stellar wind, and their impact on the two shortest-period planets, which will be presented in the second paper of this series. | \label{sec:intro} Both radial velocity and photometric transit surveys have revealed the presence of a great diversity in the structure and geometry of planetary systems, and in the physical characteristics of the planets composing them (e.g., \citealt{mullally2015}, \citealt{winn2015}). Although we have already discovered several hundreds of planetary systems, none of them resembles the solar system. Multi-planetary systems (i.e. systems with more than one planet) allow us to infer more about the formation, structure, and evolution of planets compared to single-planet systems. For example, the presence and detection of gravitational interactions between planets composing a planetary system allows one to measure their masses through radial velocity and/or transit timing variations, and possibly to detect further planets in the system \citep[e.g.,][]{ttv,cloutier2017}. In addition, the presence of mutual inclination between the orbits of planets in the same system gives us clues about their past interaction and future stability of the system \citep[e.g.,][]{veras2004}. Key parameters allowing us to infer something about the past formation and evolution history of planets and of their orbits are the (sky-projected) obliquity (i.e. the angle between the angular momentum vector of the host star and that of the planetary orbit) and the mutual inclination among the orbits in a planetary system. On the basis of the solar system, for which the obliquity is of about 6$^{\circ}$, one might expect an average good alignment between the stellar rotation and the orbits of planets. This is in fact generally true, particularly for host stars cooler than 6250\,K, however there are a number of systems that have been found to have significant mis-alignments, even among the cooler host stars \citep[e.g.,][]{queloz2010,winn2010,albrecht2012,bourrier2018}. The exact reasons for these mis-alignments are unknown and a number of theories, connected with either the star, the planets, or neighbouring stars, have been put forward to explain the observations \citep[e.g.,][]{fabrycky2007,chatterjee2008,batygin2012,rogers2012,spalding2014}. A low obliquity and an orbital alignment within a given system suggest that the system did not go through strong dynamical interactions and that migration happened quietly within the disc. In contrast, mis-aligned systems and large mutual inclinations among planets in a system are strong indications of the presence of significant planet-planet scattering through the Kozai-Lidov mechanism \citep[e.g.,][]{winn2015}. For closely packed systems of small planets, even small mutual inclinations can potentially lead to orbit crossings, thus planetary collisions \citep{veras2004}. Measuring obliquities is therefore one of the keys to unravel the history of planetary systems, but it is not straightforward as it requires the knowledge of the planetary orbital and stellar inclination angles. Planetary orbital inclination angles can be derived almost exclusively for transiting planets \citep[see][for the case of a non-transiting planet]{brogi2013} by fitting the transit light curve. Measuring stellar inclination angles is more complicated, though. The most direct method consists of comparing the stellar projected rotational velocity (\vsini), measurable from high-resolution spectra, and the stellar rotational velocity, measurable from the stellar rotation period \citep[e.g., from spot crossings observable in light curves;][]{mcquillan2014} and radius \citep[see e.g.,][]{hirano2012,hirano2014,walkowicz2013,morton2014}. Further extremely powerful methods are Doppler tomography \citep[e.g.,][]{gandolfi2012,zhou2016} and that based on the detection and characterisation of the Rossiter-McLaughlin effect, which describes the stellar rotation along the transit chord, thus giving the possibility to measure the alignment between the orbit of a planet and the stellar rotation axis \citep[e.g.,][]{cegla2016,bourrier2018}. The stellar inclination angle can also be measured by employing the Zeeman Doppler Imaging technique, employed in this work, that can be then combined with the measured orbital inclination of the transiting planets to derive the obliquity. Planets also allow us to gather precious information about their host stars, for example regarding stellar winds, that would otherwise be impossible to obtain. The winds of late-type stars are optically thin and hence extremely difficult to detect and study. Several indirect methods have placed upper limits on these winds, such as the detection of radio \citep[e.g.,][]{fichtinger2017} and X-ray emission \citep[e.g.,][]{wargelin2002}. Stellar winds can actually be detected and their properties inferred from detections of astrospheric absorption \citep[e.g.,][]{wood2004}. However, detections of stellar winds through astrospheric absorption are currently available for only a small number of objects. The observation and modelling of the interaction between the stellar wind and the extended atmosphere of a planet allows one to instead directly infer the physical properties of the wind (i.e.\ temperature and velocity) at the position of the planet \citep[e.g.,][]{bourrier2013,kislyakova2014,vidotto2017}. The ideal case is that of a multi-planet system in which the stellar wind-planetary atmosphere interaction is detected for more than one planet. This would provide constraints on the stellar wind temperature and velocity at multiple distances from the star, while also knowing that, for example, the stellar wind mass-loss rate is always the same, hence providing a solid anchor on the stellar wind density and velocity. Direct magnetic field detections for stars hosting planets are quite rare \citep[e.g.][]{Vidotto2014, Mengel2017}. Magnetic activity makes detecting planets difficult, thus successful planet searches generally are biased towards stars with the least magnetic activity. Such weak magnetic fields are difficult to detect, requiring long exposures even for very bright stars, thus very few stars have both planet detections and magnetic field measurements. However, observational knowledge of a star's magnetic field is necessary for understanding how the wind is sculpted by the magnetic field, and thus the impact of the wind on planets. Given the general lack of stellar wind detections, and paucity of magnetic field measurements for planet hosts, there are very few planet hosting stars for which we have information about both the wind and magnetic field, and until now none for which the observations were contemporaneous. The multi-planet system HD\,219134 is key to further advancing our understanding of planets and stellar winds. With five close-in planets (orbital separation $a$ smaller than 0.4\,AU) and one distant gaseous giant planet ($a$\,$\approx$\,3\,AU), HD\,219134 is one of the few systems known to date that very roughly resembles the solar system's architecture, and it lies just 6.5\,pc away from us \citep{motalebi2015,vogt2015}. The six planets orbit a 0.81\,\Mo\ main-sequence K3 star \citep[the stellar radius is 0.778\,\Ro;][]{boyajian2012,gillon2017}, with an estimated age of 11.0$\pm$2.2\,Gyr. The old age is confirmed by the long stellar rotation period of about 40 days and by the low average value of the \logR\ stellar activity parameter of about $-$5.02 \citep{motalebi2015,vogt2015}. For comparison, the basal chromospheric flux level of main-sequence late-type stars is \logR\,=\,$-$5.1 \citep{wright2004} and the average solar \logR\ value is $-$4.902$\pm$0.063 \citep[95\% confidence level;][]{mamajek2008} and ranges between a minimum of about $-$5.0 and a maximum of about $-$4.8 along the solar activity cycle. HD\,219134 has a 11.7 year chromospheric activity cycle \citep{johnson2016}, similar to the Sun. A surface average longitudinal magnetic field of 1.1$\pm$0.1\,G was detected for \hd\ \citep{marsden2014}. The planets were first detected by radial velocity and, making use of {\it Spitzer} light curves, \citet{motalebi2015} and \citet{gillon2017} discovered that the two innermost planets are transiting. This allowed a precise measurement of the planetary densities, revealing that both planets have an Earth-like density, with less than 10\% uncertainty \citep{gillon2017}. The inner most planet (b) has a mass, radius, and equilibrium temperature of $4.74 \pm 0.19$ \Me, $1.602 \pm 0.055$ \Re, and 1015 K, while the further out transiting planet (c) has a mass, radius, and equilibrium temperature of $4.36 \pm 0.22$ \Me, $1.511 \pm 0.047$ \Re, and 782 K \citep{gillon2017}. The outer planets are not transiting, due to their distance from the star, and hence their radii are unknown and the measured masses are lower limits. The rather large density of the two transiting innermost planets suggests the lack of a hydrogen-dominated envelope. This is confirmed by the low value of the restricted Jeans escape parameter $\Lambda$ \citep{fossati2017} of the two planets ($\Lambda$\,$\simeq$\,22 and 28 for \hd\,b and c, respectively), implying that a hydrogen-dominated envelope would have likely escaped within a few hundred Myr \citep{fossati2017}. \citet{DornHeng2018} arrived at the same conclusion by employing a Bayesian inference method based on the stellar properties and modelling the escape through the energy-limited approximation. \citet{tian2009} showed that super-Earths similar to \hd\,b and c, subjected to hundreds times more high-energy (XUV) stellar flux than the Earth at present, would completely lose their CO$_2$ content within about 1\,Gyr. Considering that the star was more active in the past and that the work of \citet{tian2009} is based on planets in the habitable zone ($\approx$300\,K), the two planets have likely lost most, if not all, of their secondary atmosphere. One can therefore expect that \hd\,b and c have lost both primary, hydrogen-dominated, and secondary, CO$_2$-dominated, atmospheres because of the high temperature and high-energy stellar radiation. The close proximity to the star of planets a and b, and their lack of a dense atmosphere, implies a dense stellar wind impacting on the planetary surfaces. It is expected that the surface of both planets sputter atoms and molecules, similar to what occurs on Mercury. The elements released from the planetary surface would build up a thin, metal-rich exosphere \citep{schaefer2009,ito2015}, consisting mostly of Na, O, Si, and Fe atoms/ions \citep[e.g.,][]{miguel2011,kite2016}. Some of the atoms may dissociate and ionise, and their structure and velocity would then be controlled by the stellar wind properties and the interplanetary magnetic field carried by the stellar wind. The structure of this exosphere is discussed in the second paper in this series (Vidotto et al.\ 2018b). The high-energy stellar flux and stellar wind therefore play a fundamental role in shaping the evolution of these planets, and in controlling the formation and characteristics of a metal-rich exosphere. In particular, the stellar wind drives the sputtering processes leading to the formation of a metal-rich exosphere, while the XUV flux is mostly responsible for the ionisation processes on-going in the thin exosphere. The main objective of this paper to place the strongest observational constraints currently possible on the magnetic field, wind, and XUV flux of \hd. This is the first in a series of two works aiming at constraining the properties of the stellar wind and XUV flux, and using this to model the planetary exospheres. In this first paper, we present the results of spectropolarimetric, spectroscopic, and photometric observations. We derive the map of the surface magnetic field and analyse the stellar Ly$\alpha$ line to constrain the wind mass-loss rate. We further analyse space-based ultraviolet (UV) observations obtained with {\it AstroSat}. The second paper will then be dedicated to modelling the stellar wind and planetary metal-rich exosphere. This work is structured as follows. In Sect.~\ref{sec:specparams} we present the spectropolarimetric observations of \hd\ and derive its stellar parameters, and in Sect.~\ref{sec:zdi} we present the results of the Zeeman Doppler Imaging analysis. In Sect.~\ref{sec:astrosphere} we present the analysis of the stellar Ly$\alpha$ emission line of \hd\ and the astrospheric detection. In Sect.~\ref{sec:astrosat} we show the {\it AstroSat} observations, around the transit of \hd\ b, and the resulting upper limit on the detectability of planet b at near-ultraviolet wavelengths. In Sect.~\ref{sec:xuv}, we derive the high-energy fluxes of the star and high-energy flux estimates at the distance of planets b and c. In Sect.~\ref{sec:conclusions} we discuss the results of the observations and gather the conclusions. | 18 | 8 | 1808.00406 |
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1808 | 1808.07010_arXiv.txt | We have obtained $> 10$ hours of medium resolution ($R \sim 15000$) spectroscopic exposures on the transiting exoplanet host star WASP-12, including $\sim2$ hours while its planet, WASP-12b, is in transit, with the Hobby-Eberly Telescope (HET). The out-of-transit and in-transit spectra are coadded into master out-of-transit and in-transit spectra, from which we create a master transmission spectrum. Strong, statistically significant absorption features are seen in the transmission spectrum at H$\alpha$ and \ion{Na}{1} (the Na D doublet). There is the suggestion of pre- and post-transit absorption in both H$\alpha$ and \ion{Na}{1} when the transmission spectrum is examined as a function of phase. The timing of the pre-transit absorption is roughly consistent with previous results for metal absorption in WASP-12b, and the level of the \ion{Na}{1} absorption is consistent with a previous tentative detection. No absorption is seen in the control line of \ion{Ca}{1} at $\lambda$6122. We discuss in particular whether or not the WASP-12b H$\alpha$ absorption signal is of circumplanetary origin---an interpretation that is bolstered by the pre- and post-transit evidence---which would make it one of only a small number of detections of circumplanetary H$\alpha$ absorption in an exoplanet to date, the most well-studied being HD 189733b. We further discuss the notable differences between the HD 189733 and WASP-12 systems, and the implications for a physical understanding of the origin of the absorption. | \label{s:intro} \subsection{The Hot Jupiter WASP-12b} \label{ss:wasp12b} WASP-12b is a very short-period hot Jupiter \citep{Hebb2009} discovered with the SuperWASP program \citep{Pollacco2006}. At the time of its discovery it was the most highly-irradiated and shortest period planet in the literature, and has been the subject of significant attention ever since. Shortly after its discovery, \citet{Fossati2010} found evidence suggestive of absorption by \ion{Mg}{2} and other metals in WASP-12b's atmosphere through {\it HST}/COS observations. These results include an indication of pre-transit absorption, potentially signifying that WASP-12b has a significant extended atmosphere. \citet{Haswell2012} presented a second {\it HST}/COS visit that bolstered the results of \citet{Fossati2010}, including a detection of \ion{Fe}{2}. The \citet{Haswell2012} results showed that the required interpretation of the system was more complex, especially given that the ingress of the second visit began significantly earlier than the first visit in \citet{Fossati2010}. The depth of the transits indicate that the planet's Roche lobe is overfilled \citep{Lai2010,Debrecht2018}, and there is also observational and theoretical evidence that material from WASP-12b forms a torus- or disk-like structure around the star \citep{Haswell2012,Fossati2013,Debrecht2018}. The C/O ratio and water content of WASP-12b have also been the focus of much study, especially the possibility that the planet is carbon-rich rather than oxygen-rich. Evidence of a high C/O ratio and a weak thermal inversion was found by \citet{Madhusudhan2011}. More recently, \citet{Kreidberg2015} found evidence for water in the HST/WFC3 transmission spectrum of WASP-12b. This detection is potentially consistent with high C/O ratios ($>1$) under certain assumptions, but strongly favors a C/O ratio closer to 0.5. Another question of interest is whether WASP-12b has an absorbing layer of TiO/VO. The nominal expectation from using the classification system of \citet{Fortney2008} is that WASP-12b should have TiO/VO in its atmosphere that masks \ion{Na}{1} absorption. However, \citet{Sing2013} found evidence of a lack of TiO in WASP-12b in {\it HST} observations. Recently, \citet{Burton2015} tentatively detected \ion{Na}{1} in WASP-12b's atmosphere through defocused transmission spectroscopy at the level of $0.12\pm0.03[\pm0.3]$\%, with the error in brackets representing additional systematic uncertainty; accounting for the uncertainty they state that $0.15$\% is a better representation of their absorption measurement. This detection of \ion{Na}{1} is consistent with the \citet{Sing2013} results and further indicates that WASP-12b does not fit neatly into the \citet{Fortney2008} classification system. Furthermore, \citet{Sing2013} find that their transmission spectrum can be fit equally well by either Rayleigh or Mie scattering. However, Mie scattering is a much better fit to the expected atmospheric temperatures (and the observed blackbody emission spectrum), and implies a high-altitude haze. This subsequently means that any observed line absorption must occur at high altitudes above the haze. One final recent WASP-12b observation is relevant to the current work, a search for \ion{He}{1} absorption at 10833 \AA{} by \citet{KriedbergOklopcic2018}. They found a transit depth of $59\pm143\ppm$ relative to the adjacent wavelength bands in Hubble Space Telescope/Wide Field Camera 3 G102 grism data. This non-detection does constrain certain models for WASP-12b's atmosphere and indicates that any helium absorption in WASP-12b is smaller than a recent detection in WASP-107b \citep{Spake2018}. However, the measurement is made over an integration band of 70 \AA{} (the instrumental resolution) and does not rule out the possibility that significant absorption might be observed with a higher instrumental resolution. \subsection{Observations of Hydrogen in Exoplanets} \label{ss:hinexo} Observations of broad hydrogen envelopes measured in ground state absorption (specifically in Ly$\alpha$) have been made in multiple planets, including the well-studied planets HD 209458b {\citep{VidalMadjar2003,VidalMadjar2004} and HD 189733b \citep{LecavelierDesEtangs2010,LecavelierDesEtangs2012}. Such observations give insight into the possibility of star-planet interactions and atmospheric escape. However, exoplanetary Ly$\alpha$ measurements pose significant observational difficulties. First, Ly$\alpha$ is in the UV, and UV-capable facilities with adequate throughput and spectral resolution are a very limited resource; only {\it HST} is capable of such observations in Ly$\alpha$ at the current time. Second, certain stars are not bright at Ly$\alpha$, especially sun-like stars. Furthermore, interstellar Ly$\alpha$ absorption is significant, and absorbs stellar Ly$\alpha$; this absorption poses difficulties for all observations of exoplanetary Ly$\alpha$ and makes them completely impractical in lines of sight longer than 50 pc or so. In short, only certain relatively nearby systems will have adequate observable Ly$\alpha$ flux available to perform exoplanetary transmission spectroscopy. The aforementioned detection of \ion{He}{1} at 10833 \AA{} in WASP-107b \citep{Spake2018} is one promising observational strategy for observing extended exoplanetary atmospheres without UV facilities. Another option is to directly observe hydrogen, albeit not in the ground state, through its Balmer series of transitions at visible wavelengths. In a series of papers \citep[][hereafter Papers I, II, and III]{Redfield2008, Jensen2011, Jensen2012} the current authors examined the Hobby-Eberly Telescope (HET) transmission spectra of four hot Jupiter-like planets for absorption due to \ion{Na}{1}, \ion{K}{1}, and H$\alpha$. Paper I made the first ground-based observation of an exoplanetary atmosphere by detecting \ion{Na}{1} in HD 189733b. Paper II confirmed previously observed atmospheric \ion{Na}{1} absorption in HD 189733b (Paper I) and HD 209458b \citep{Charbonneau2002}, and found the hint of possible \ion{Na}{1} absorption in HD 149026b. In Paper III, H$\alpha$ was detected in HD 189733b's transmission spectrum, the first-ever such detection of exoplanetary H$\alpha$. Because the HET observations do not lend themselves to complete light curves that encompass an entire transit, observations with Keck I/HIRESr were obtained to observe a full transit of HD 189733b. The first observations, in 2013, confirmed the H$\alpha$ transit absorption and found corresponding H$\beta$ and H$\gamma$ transit absorption, along with a significant pre-transit signal in all three lines \citep{Cauley2015}. Subsequent observations in 2015 found strong variation in the pre-transit and transit signals, leading to an uncertain physical model for the geometry of the absorption \citep{Cauley2016}. \citet{Barnes2016} challenged the circumplanetary interpretation of the H$\alpha$ absorption observed in HD 189733b, presenting an alternate interpretation that the transmission spectrum is dominated by contrast effects during transit. Any star's spectrum will vary over its disk (due to spots or other active regions), and it is possible that a transiting planetary disk may result in differential spectroscopy that mimics excess absorption in certain lines. However, there remain two significant arguments for a circumplanetary origin of the H$\alpha$ absorption in HD 189733b. First, short-cadence monitoring of HD 189733's stellar H$\alpha$ shows that variation at the level of the observed pre-transit signals (which cannot be explained by the contrast effect) is uncommon \citep{Cauley2017b}. Second, modeling of the stellar contrast effect that might mimic absorption during transit demonstrates that creating the absorption signal seen in HD 189733b requires a highly constrained distribution of active regions distributed only along the transit chord and/or a potentially implausible level of stellar activity for HD 189733 \citep{Cauley2017c}. As of early 2018, HD 189733b was the only exoplanet with a published detection of H$\alpha$ absorption. However, there are two recent detections of H$\alpha$ in exoplanetary atmospheres, both in planets around A stars---KELT-9b \citep{Yan2018} and MASCARA-2b/KELT-20b \citep{Casasayas2018}. \citet{Yan2018} detected 1.15\% of extra absorption at H$\alpha$ line center in KELT-9b, and attribute the H$\alpha$ absorption to a hot extended atmosphere that is driven by the intense UV radiation of the star. In addition to finding H$\alpha$, \citet{Casasayas2018} detected \ion{Na}{1} in KELT-20b. They found that a temperature higher than equilibrium (4210 K vs.~2260 K, respectively) is required in order to explain the observations, which could be explained by the large amount of UV energy delivered by the central star. \subsection{WASP-12b and Comparative Planetology} \label{ss:planetology} In this paper we present the HET transmission spectrum of WASP-12b. This target was selected for HET observations as a point of comparison to the Paper III HD 189733b H$\alpha$ detection \citep[and prior to the Keck follow-up in][]{Cauley2015, Cauley2016}. The central star WASP-12 has been classified as a G0V star by \citet{Bergfors2013}; however, \citet{Fossati2010star} find a temperature of $6250\pm100$ K, consistent with the value \citet{Hebb2009} derive, which would suggest a spectral type of approximately F7. In any case, as it pertains to the possibility of H$\alpha$ absorption, WASP-12b is closer to a hotter central star than HD 189733b, and thus more highly irradiated. On the other hand, HD 189733 is a later-type (K0V) star that is presumably more active than WASP-12, and its Ly$\alpha$ emission is likely to be both stronger and more variable. This is significant as \citet{Huang2017} modeled the H$\alpha$ emission in HD 189733b, and found that stellar Ly$\alpha$ emission and Lyman continuum emission play an important role in creating the $n=2$ hydrogen population, although collisional excitation also plays a role \citep{Christie2013}. Furthermore, \citet{Huang2017} suggest that HD 189733's activity may be the cause of the variation in the H$\alpha$ transit depth between \citet{Cauley2015} and \citet{Cauley2016}. Thus, any H$\alpha$ detected in WASP-12b, given its intermediate spectral type and different physical conditions, would immediately be an interesting data point to enable comparative planetology for any exoplanets exhibiting transit-correlated H$\alpha$ absorption, including HD 189733b, KELT-9b, and KELT-20b. Table \ref{table:parameters} provides information on the WASP-12 system. \begin{deluxetable*}{ccc} \tablecolumns{3} \tabletypesize{\small} \tablecaption{WASP-12 System Parameters\tablenotemark{a}\label{table:parameters}} \tablehead{\colhead{Parameter} & \colhead{WASP-12/WASP-12b} & \colhead{Unit}} \startdata Transit Midpoint & $2454508.97682\pm0.0002$ & HJD \\ Period & $1.09142245\pm3\times10^{-7}$ & days \\ Transit Duration & $0.122\pm0.001$ & days \\ $R_P$ & $1.79\pm0.09$ & $R_{\rm Jupiter}$ \\ $R_\star$ & $1.63\pm0.08$ & $R_\odot$ \\ $a/R_\star$ & $2.98\pm0.154$ & N/A \\ $b$ & $0.375^{+0.042}_{-0.049}$ & N/A \\ $i$ & $82.5^{+0.8}_{-0.7}$ & degrees \\ \enddata \tablenotetext{\rm a}{All values from \citet{Maciejewski2011} and references therein, including \citet{Hebb2009}.} \end{deluxetable*} In \S\ref{s:obsdata} we describe our observations and data reduction, and in \S\ref{s:analysis} we describe our analysis methods. Our results are presented in \S\ref{s:results}. We discuss our results and possible future work in \S\ref{s:conclusions}. | \label{s:conclusions} \subsection{Discussion} \label{ss:discussion} As we have stated, detection of H$\alpha$ in exoplanetary atmospheres is rare. The H$\alpha$ in HD 189733b \citep[Paper III,][]{Cauley2015,Cauley2016} is by far the best studied. There are also a handful of non-detections of H$\alpha$ such as for HD 209458b in \citet{Winn2004} and HD 147506b and HD 149026b in Paper III. Recent additional detections in KELT-20b \citep{Casasayas2018} and KELT-9b \citep{Yan2018} are intriguing, especially considering that they are around A stars, but they are not as well-studied as HD 189733b. \citet{Salz2016} modeled the atmosphere of several hot Jupiter planets and found a mass-loss rate for WASP-12b that is large (3.4\% per Gyr) but significantly smaller than previous results \citep{Li2010,Lai2010}, thus resolving the prior implication that the planet was short-lived and its discovery statistically improbable. The mass-loss rate found by \citet{Salz2016} indicates a very large hydrogen envelope, and a large theoretical Ly$\alpha$ signal that absorbs over 80\% of the incoming flux at line center and spans hundreds of $\kmpers$. This model includes absorption that is significant well beyond the Roche lobe. Given the abundance of hydrogen, the detection of H$\alpha$ is not particularly surprising, nor is the fact that we detect pre- and possibly post-transit absorption. The primary open question remains what mechanisms generate $n=2$ in significant amounts, and whether these mechanisms can explain the observational details, particularly the depth and velocity range, of the H$\alpha$ absorption. What creates $n=2$ hydrogen in hot Jupiter atmospheres? The HD 189733b H$\alpha$ detection has been modeled in detail by \citet{Christie2013} and \citet{Huang2017}. The former used a hydrostatic atmosphere model, but the Ly$\alpha$ radiation was not modeled in detail. \citet{Christie2013} found that a reasonably constant $n=2$ hydrogen density within the atomic layer could be created if collisional excitation dominated. \citet{Huang2017} expanded on the work of \citet{Christie2013} by including a more detailed treatment of the Ly$\alpha$ radiative transfer. By considering Ly$\alpha$ coming from recombinations within the atmosphere, the radiative excitation rate can exceed the collisional excitation rate, meaning that the excitation can occur at significant levels deeper within the atmosphere at greater $n=1$ hydrogen densities. Given that WASP-12's Ly$\alpha$ emission is presumably weaker than HD 189733's, it is unclear how significant $n=2$ could exist. However, WASP-12b's atmosphere should be more extended than HD 189733b's, with a larger scale height. WASP-12b's lower atmosphere is certainly hotter than HD 189733b's, though this may not necessarily be the case at higher altitudes \citep{Salz2016}. If and where WASP-12b's atmosphere is hotter, it is possible that even with weaker EUV radiation, collisions will contribute significantly to $1s\rightarrow2s$ excitation. Our results also must be understood in the context of the evidence for a torus- or disk-like structure around WASP-12, made up of material from WASP-12b \citep{Lai2010,Haswell2012,Fossati2013,Debrecht2018}. We show a clear correlation with transit for H$\alpha$ and \ion{Na}{1}, which argues for some asymmetry in the disk, such as a collision between the accretion stream from the planet and the disk \citep{Lai2010}, which could be the source of early ingress. Understanding the early ingress also has significant implications as a potential probe of the magnetic field \citep{Vidotto2010,Llama2011}. The evidence for late egress in our observations is not as compelling as the evidence for early ingress, but it is also worth further study. In \S\ref{ss:hinexo} we noted that there are observational challenges to relying on Ly$\alpha$ as a diagnostic of extended exoplanetary atmospheres. First, UV instrumentation that covers Ly$\alpha$ at adequate S/N and resolution is essentially restricted to {\it HST} at the current time. Second, the stellar flux at Ly$\alpha$ varies significantly as a function of spectral type. Active, late-type stars have significant line emission at Ly$\alpha$; this includes K stars like HD 189733b. Late F stars like WASP-12 have line emission at Ly$\alpha$ but it is weaker (relative to the star's continuum) than Ly$\alpha$ emission for K or M stars, while earlier type stars like the A stars KELT-9b and KELT-20b will have significant UV continuum flux. Third, even with adequate UV instrumentation and stellar flux at Ly$\alpha$, the observation of Ly$\alpha$ transit absorption is complicated by interstellar absorption and airglow in the Earth's atmosphere, issues that are not present for H$\alpha$. Given that H$\alpha$ also probes the structure of a planet's upper atmosphere and its location in the visible red portion of the spectrum makes it accessible from the ground, H$\alpha$ is an arguably underutilized diagnostic for observing extended atmospheres. Specifically, while not all exoplanets with observable Ly$\alpha$ signatures will have corresponding, observable H$\alpha$ absorption (e.g., HD 209458b), using both can be complementary approaches insofar as Ly$\alpha$ observations are flux limited. This is the case with WASP-12b; at a distance of approximately 430 pc \citep{Gaia2018}, interstellar absorption prevents Ly$\alpha$ observations of WASP-12 with {\it HST}. The recent \ion{He}{1} detection by \citet{Spake2018} is another way in which extended atmospheres may potentially be observed without UV observations of Ly$\alpha$. Furthermore, the far- and near-UV metals line observations discussed in \S\ref{ss:wasp12b} represent another significant way in which extended atmospheres may be probed. In addition to being in different wavebands, all of these methods represent at least somewhat different diagnostics of exoplanetary atmospheres, and therefore they are all useful as different probes of exoplanetary characteristics. \subsection{Summary} \label{ss:summary} We have presented the transmission spectrum of WASP-12b from HET observations in March/April 2012. The spectrum shows clear features at H$\alpha$ (Fig.~\ref{fig:Hatrans}) and \ion{Na}{1} (Fig.~\ref{fig:Natrans}) while no obvious features are observed at \ion{Ca}{1} (Fig.~\ref{fig:Catrans}, intended as a control line) or H$\beta$ (Fig.~\ref{fig:Hbtrans}). The H$\alpha$ absorption marks only the fourth such detection in an exoplanetary atmosphere after HD 189733b, KELT-20b \citet{Casasayas2018}, and KELT-9b \citep{Yan2018}, while the \ion{Na}{1} absorption is roughly consistent with previous results by \citet{Burton2015} and likewise one of larger, but still limited, number of exoplanetary \ion{Na}{1} detections. Phase curves of H$\alpha$ (Fig.~\ref{fig:HaplhaPhase09}) and \ion{Na}{1} (Fig.~\ref{fig:NIPhase}) indicate the possibility of pre- and post-transit absorption, roughly consistent with the results of \citet{Fossati2010} and \citet{Haswell2012} for metals in the extended atmosphere of WASP-12b. The evidence for pre- and post-transit absorption suggests that stellar effects such as the contrast effect are not an adequate explanation for our observations. The lack of H$\beta$ absorption at a level that would be consistent with the strong H$\alpha$ signal is somewhat puzzling and the biggest challenge to our interpretation of the H$\alpha$ transmission spectrum as a real, astrophysical signal from WASP-12b's circumplanetary material; however, our results for the H$\alpha$/H$\beta$ ratio are not ruled out to high statistical significance in the optically thin case. \subsection{Future Work} \label{ss:future} As with HD 189733b, the transit-correlated H$\alpha$ absorption cannot be fully understood with our track-limited observations made by the HET. The H$\alpha$ profile of WASP-12b needs to be observed at high S/N over an entire single transit in order to get a better picture of what the correlation between transit and absorption actually is. Our previous work has shown that absorption detected through the HET can be well-characterized through a large telescope that observes continuously through transit, e.g., with Keck \citep{Cauley2015, Cauley2016}. Such observations would also allow us to more rigorously characterize the suggestion of pre- and post-transit absorption in our observations, as was noted in HD 189733b by \citet{Cauley2015,Cauley2016}. | 18 | 8 | 1808.07010 |
1808 | 1808.02897_arXiv.txt | { It is generally assumed that for energy injection before recombination, all of the injected energy is dissipated as heat in the baryon-photon plasma, giving rise to the $y$-type, $i$-type, and $\mu$-type distortions in the CMB spectrum. We show that this assumption is incorrect when the energy is injected in the form of energetic (i.e. energy much greater than the background CMB temperature) particles. We evolve the electromagnetic cascades, from the injection of high energy particles, in the expanding Universe and follow the non-thermal component of \normalcolor CMB spectral distortions resulting from the interaction of the electromagnetic shower with the background photons, electrons, and ions. The electromagnetic shower loses a substantial fraction of its energy to the CMB spectral distortions before the energy of the particles in the shower has degraded to low enough energies that they can thermalize with the background plasma. This spectral distortion is the result of the interaction of non-thermal energetic electrons in the shower with the CMB and thus has a shape that is substantially different from the $y$-type or $i$-type distortions. The shape of the final \emph{non-thermal relativistic} ($ntr$-type) CMB spectral distortion depends upon the initial energy spectrum of the injected electrons, positrons, and photons and thus has information about the energy injection mechanism e.g. the decay or annihilation channel of the decaying or annihilating dark matter particles. The shape of the spectral distortion is also sensitive to the redshift of energy injection. Our calculations open up a new window into the energy injection at $z\lesssim 2\times 10^5$ which is not degenerate with, and can be distinguished from the low redshift thermal $y$-type distortions.} \notoc \begin{document} | 18 | 8 | 1808.02897 |
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1808 | 1808.00540_arXiv.txt | We discuss the use of optical cavities as tools to search for dark matter (DM) composed of virialized ultra-light fields (VULFs). Such fields could lead to oscillating fundamental constants, resulting in oscillations of the length of rigid bodies. We propose searching for these effects via differential strain measurement of rigid and suspended-mirror cavities. We estimate that more than two orders of magnitude of unexplored phase space for VULF DM couplings can be probed at VULF Compton frequencies in the audible range of 0.1-10 kHz. | 18 | 8 | 1808.00540 |
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1808 | 1808.05962_arXiv.txt | Galaxies within galaxy clusters are known to be subject to a wide variety of environmental effects, both gravitational and hydrodynamical. In this study, we examine the purely gravitational interaction of idealised galaxy models falling into a galaxy cluster in the context of Modified Newtonian Dynamics (MOND). This modification of gravity gives rise to an external field effect (EFE), where the internal dynamics of a system are affected by the presence of external gravitational fields. We examine the consequences of the EFE on low and high mass disk galaxies in time-evolving analytic background cluster potentials, considering orbits with weak and strong tidal fields. By varying the orbital plane of the galaxies we also test the effect of having the tidal interaction orthogonal or parallel to the disk. Furthermore, we consider as a control sample models where the EFE has been removed and they are only affected by tides. Our results suggest that MOND cluster galaxies should exhibit clear asymmetries in their isophotes, suffer increased mass loss and a reduction in their rotation curves due to the combined effect of cluster tides and the external field. In particular, low mass galaxies are hit hard by the EFE, becoming dominated by dispersion rather than rotation even in the absence of tides. | The ``missing mass'' problem is one of the classic problems of astronomy. In the context of individual disk galaxies, this missing mass, referred to as \textit{dark matter}, manifests itself in the non-Keplerian behaviour of their rotation curves. This additional unseen mass is also required to explain observations of the velocity dispersions of galaxies within galaxy clusters as well as the X-ray temperature profile of the cluster gas. At the level of cosmology, dark matter is required to explain the power spectrum of temperature fluctuations in the CMB, and the process of structure formation more generally. Direct and indirect detection searches for dark matter particles have made impressive improvements in sensitivity over the years, yet an unambiguous signal has so far proven elusive (see \citealp{DMpaper} for a recent review). Furthermore, from an aesthetic point-of-view, it may seem unsatisfactory to introduce a new particle species, with properties constructed precisely to elude immediate detection. Combined with the ongoing mystery of dark energy, the current standard cosmological model may be considered as a model with a (small) number of free parameters which can be adjusted to match observations extremely well, particularly at high redshift. The freedom within this model represents the entirely unknown nature of the dark sector: effectively a parameterisation of our ignorance. Furthermore, well into the non-linear regime, there are potential small-scale problems with the $\Lambda$CDM model \citep{smallscale}, although these may eventually be resolved by improved baryonic physics in simulations \citep{baryonicPhysSim}. In addition, the recently discovered planes of satellites around the Milky Way, Andromeda and now Centaurus A \citep{centA}, present a potentially significant challenge to the standard picture of cold dark matter (for a review, see \citealp{pawlowskiPlanes}). Given these issues and the fact that the evidence for the existence of dark matter is (so far) entirely derived from gravitational dynamics, it remains an intriguing possibility to consider a modification of gravity as being responsible for the missing mass phenomenology. A very well-known and studied modification is that of Milgrom's Modified Newtonian Dynamics \citep{milgrom1,milgrom2,milgrom3}, referred to as MOND (for a thorough review see \citealp{mondreview}). The most common formulation of this paradigm is that of a modified gravitational potential, arising from an adjustment of the standard Poisson equation of Newtonian gravity. This adjustment is a function of the local accelerations, such that, in circumstances of weak accelerations below an empirically determined scale, the gradient of the gravitational potential is increased, and thus gravity is ``strengthened.'' Such a modification of gravity has interesting consequences throughout all of stellar and galactic dynamics \citep{kroupa}, suggesting the need to study numerical N-body simulations within this regime. Several such codes are now available, such as PoR \citep{por}, N-MODY \citep{nmody} and a MOND version of the cosmological code AMIGA \citep{amiga_mond}. In this work we will use the MOND N-body/hydrodynamics code known as RAyMOND \citep{raymond}, which is based upon the RAMSES code \citep{ramses}. This code has previously been used to study possible signals of MOND gravity in cosmological structure formation \citep{raymondCosmo}. In this study we work at the smaller scale of idealised galaxy models within a galaxy cluster. One of the most spectacular examples of the effects of gravity on galactic dynamics is the tidal disruption of galaxies when they encounter the external gravitational fields of other galaxies or galaxy clusters. These gravitational interactions give rise to a wide range of environmental effects such as mergers, harassment (high-speed encounters between galaxies) and tidal stripping. A huge number of examples of such systems have been observed over the years, making clear that these effects are a crucial ingredient in the process of galaxy evolution, and can even lead to their formation, in the case of tidal dwarf galaxies. Furthermore, the tidal disruption of galaxies by the cluster potential itself is thought to be a crucial ingredient in the build-up of the intracluster light, which may require substantial stripping of massive galaxies \citep{ICLbuildup}. The central importance of these gravitational encounters for galaxy evolution compels us to consider them in the context of the MOND paradigm, possibly as a means to distinguish between MOND and dark matter models, as well as to explore the process of galaxy evolution in such modified gravity theories. The main aim of our study is to examine the consequences of the external field effect on disk galaxies as they fall into a cluster. Previous analytical studies of the EFE have focussed on warping of galactic disks \citep{brada_warps} and the distortions of isophotes in giant ellipticals well within galaxy clusters \citep{elliptical_efeI}. A numerical approach (without full N-body dynamics) was taken to analyse the consequences of the EFE on the velocity dispersions of satellite dwarfs of the Milky Way \citep{efe_mw}. Beyond this, there have been a small number of studies of the dynamical effect of the EFE using N-body simulations. The study of \cite{elliptical_efeII} searched for equilibrium N-body models of elliptical galaxies within an external field, while the studies of \cite{tidal_streamsI,tidal_streamsII} analysed the evolution of globular clusters in the external field of a host galaxy, finding that the EFE both distorts the shape of the globular clusters and induces asymmetries in the tidal streams. Recent analytical work has discussed the implications of ultra-diffuse galaxies within clusters for MOND \citep{milgromUDG}. In this paper, however, we will analyse the effect of the external field on isolated galaxies that fall into a galaxy cluster, thus accounting for a time-varying EFE by including the infall process. One would expect the strength of the EFE in a low mass galaxy as it falls into a massive cluster to be considerable, given the large mass ratio and associated large difference in external and internal accelerations, and therefore potentially observable. We begin with a brief description of the formulation of MOND used in our study in Section~\ref{sec:theory}, before the simulations are discussed in Section~\ref{sec:sims}. Our results are presented in Section~\ref{sec:results} and finally our conclusions are in Section~\ref{sec:conclusions}. | \label{sec:conclusions} We have studied idealised models of low and high mass MOND disk galaxies falling into galaxy clusters of different masses. For more control over the orbital trajectories, we have used an analytic NFW model for the background cluster, with a time varying mass in order to approximate the mass growth of the clusters and ensure bound infall trajectories. Our models do not incorporate dynamical friction, although this is only significant for the most massive galaxies in galaxy clusters. The orbits of the galaxy models are either perpendicular to the disk, or parallel, with either a retrograde or prograde motion around the cluster centre in the latter case. Finally, we have analysed the consequences of removing the MOND external field effect. The main conclusion of our study is that both the low and high mass MOND galaxies are substantially distorted by the EFE as they fall into the cluster, leading to significant asymmetries in the galactic disks, even in the absence of any tidal interaction. In the presence of such an interaction, however, these asymmetries are further enhanced. The weakened potentials of the MOND galaxies within the cluster also leads to mass loss in the absence of tides and increased susceptibility to mass loss due to tidal stripping, loss of their rotational components and a morphological evolution towards a more spherical distribution. More specifically, our study has found that: \begin{itemize} \item MOND disk galaxies in clusters have distorted isophotal contours due to the EFE, with low mass galaxies in a high mass cluster the most severely affected. When further combined with a tidal interaction, the distortion is even more enhanced, producing highly asymmetric forms. The effect is weaker for high mass galaxies, but still long-lived (for all models the asymmetries last at least $1$ Gyr, and up to $4$ Gyr for the low mass models) and likely to be clearly detectable. \item All of our model galaxies that are subjected to an EFE suffer a gradual mass loss even before reaching pericentre on their orbits and are more easily tidally stripped, both due to the EFE. Again, as the EFE is stronger for low mass galaxies (given a fixed background field), these results are very clear for those models. Thus our galaxies suffer more mass loss than would be expected in the absence of the EFE or for a Newtonian galaxy embedded in a dark matter halo. \item Given the modifed internal potential of the galaxies, as well as the induced disequilibrium due to the rapidly-evolving external field (effectively a form of violent relaxation), they undergo a kinematical evolution within the cluster towards being more dispersion dominated objects. Several of the low mass galaxies in the high mass cluster evolve to dispersion-supported systems in the Coma-mass cluster, regardless of tides, although they remain out of equilibrium by the end of our simulation, as indicated by their evolving morphology. \item Finally, this loss of rotational dominance is reflected in a morphological change from disks to more spherical distributions. This morphological change has a clear dependency on the orientation of the galactic disk relative to the orbital trajectory within the cluster, given the directional nature of the orbital perturbations induced by the tidal disruption and the EFE. \end{itemize} Furthermore, from our results we can infer that, in the absence of an external field effect, MOND galaxies within galaxy clusters would behave similarly to Newtonian galaxies within dark matter halos, at least on first infall within the cluster. Given that our models do not undergo more than one pericentric passage, it is not clear if the tidal stripping of a MOND galaxy in the absence of an EFE would be comparable to that of a Newtonian model that had already lost a significant fraction of its halo after the first pericentric passage. Given our asymmetry results, it would be of great interest to explore the prevalence of galactic asymmetries in massive nearby clusters and to compare this with values in the field. Typically this parameter is taken to be indicative of galaxy mergers in the standard context of galaxy evolution, yet in our study we show large increases in asymmetry as galaxies move within the cluster, \textit{in the absence of any galaxy-galaxy interaction}. In \cite{homeier} the asymmetry of cluster and field galaxies is compared for late-type galaxies at $z \sim 1$, where the galaxies in lower mass (X-ray faint) clusters were found to be more asymmetric than those in higher mass (X-ray luminous) clusters, a property attributed to the likelihood of more mergers in low mass clusters. In our study we suggest that \textit{high mass} clusters should have a larger population of asymmetric galaxies. Our simulations do not include the possibility of galaxy-galaxy interactions, but it is worth mentioning that mergers are expected to be much rarer in MOND \cite{kroupa} due to the lack of dynamical friction. Thus it seems likely that in MOND galaxy asymmetries are indeed more common in high mass clusters, whereas in $\Lambda$CDM they would be more common in low mass clusters/groups. A natural extension of this study, therefore, would be to test this hypothesis by considering a realistic live cluster evolution of multiple infalling galaxies and groups, allowing for possible galaxy-galaxy interactions. This would allow us to investigate the effect of dynamical friction on the galaxy interactions and on their orbital trajectories, given the clear differences expected between $\Lambda$CDM and MOND in this area \cite{SanchezSalcedoDF,NipotiDF,HongShengDF}. In addition, such a study would allow a phase-space analysis of the EFE as an environmental effect \citep{oman,jaffe1,jaffe2,rhee}. All known MOND theories based on the modification of the Newtonian Poisson equation include an external field effect. More specifically, the relativistic extensions of MOND that invoke additional degrees of freedom in the gravitational sector give rise to violations of the Strong Equivalence Principle and an external field effect. Furthermore, the recent discovery of an electromagnetic counterpart to a gravitational wave detection \citep{GravWaveDiscovery} has placed very stringent bounds on the (low redshift) propagation velocity of gravitational waves, ruling out large classes of relativistic modified gravity theories that predict gravitational wave propagation at velocities other than the speed of light \citep{GWchallenge}. If observational data of galaxies within clusters consistently contradicts the results of this study, this may present a significant challenge to the presence of an external field effect, putting pressure on the standard non-relativistic formulations of MOND. Other formulations that do not predict an external field effect (e.g. modified inertia) are known, but are much less studied, and have highly unusual theoretical properties, such as extreme non-locality \citep{modInertia1,modInertia2}. The consequences for the dynamics of galaxies and galaxy clusters of modifying gravity are profound and varied. Further investigations along these lines, using numerical simulations within the MOND context, may allow us to find further interesting and potentially observable signatures. | 18 | 8 | 1808.05962 |
1808 | 1808.08247_arXiv.txt | The Laser Interferometer Space Antenna (LISA) gravitational-wave (GW) observatory will be limited in its ability to detect mergers of binary black holes (BBHs) in the stellar-mass range. A future ground-based detector network, meanwhile, will achieve by the LISA launch date a sensitivity that ensures complete detection of all mergers within a volume $>\!\mathcal{O}(10)\,{\rm Gpc}^{3}$. We propose a method to use the information from the ground to revisit the LISA data in search for sub-threshold events. By discarding spurious triggers that do not overlap with the ground-based catalogue, we show that the signal-to-noise threshold $\rho_{\rm LISA}$ employed in LISA can be significantly lowered, greatly boosting the detection rate. The efficiency of this method depends predominantly on the rate of false-alarm increase when the threshold is lowered and on the uncertainty in the parameter estimation for the LISA events. As an example, we demonstrate that while all current LIGO BBH-merger detections would have evaded detection by LISA when employing a standard $\rho_{\rm LISA}=8$ threshold, this method will allow us to easily (possibly) detect an event similar to GW150914 (GW170814) in LISA. Overall, we estimate that the total rate of stellar-mass BBH mergers detected by LISA can be boosted by a factor $\sim\!4$ ($\gtrsim\!8$) under conservative (optimistic) assumptions. This will enable new tests using multi-band GW observations, significantly aided by the greatly increased lever arm in frequency. | 18 | 8 | 1808.08247 |
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