adsabs/WIESP2022-NER huggingface dataset:
Paper: 2020MNRAS.491.2460F (fulltext)
and the quasi-parallel case for the acceleration of the electrons could better reproduce the observed distribution of the Stokes parameter Q . For another remnant G1.90.1, which also has two opposite protrusions on the X-ray morphology, assuming it has evolved inside an elliptical planetary nebula with ears and clumps, Tsebrenko Soker ( 2015 ) investigated the formation of the shape with two ear-like protrusions using 3D HD simulations, and the interaction of the ejecta with the clumps produces the double-shock structure. SNRs evolved in inhomogeneous mediums have morphologies with interesting signatures. In this paper, we intend to study the formation mechanism of the peculiar structures illustrated by the multiband observations for SN 1006. By introducing an ambient medium with a density discontinuity, we investigate how the substructures of it are produced via numerical simulations. Different from Schneiter etal. ( 2010 ), the HD simulations in this paper are performed in 3D, and the NE and SW protrusions and the NW filament can be reproduced using the model. The model and numerical set-up for SN 1006 are shown in Section 2 . The results are presented in Section 3 . Finally, the discussion and some conclusions are given in Section 4 . 2 THE MODEL AND NUMERICAL SET-UP To investigate the dynamical evolution of the remnant SN 1006 to reproduce the peculiar morphology, we assume the ejecta have evolved in the non-uniform medium with a density discontinuity. As illustrated in Fig. 1 , the plane of the density discontinuity is perpendicular to the z -axis, and the number density of the ambient medium is assumed to be where is the unit vector along the z -axis, n H is the hydrogen number density, 1.4 is the mean atomic mass for a gas of 10:1 H:He ratio, the ratio x 3 is adopted in this paper, which is consistent with that derived from the observations with XMM Newton (Acero, Ballet Decourchelle 2007 ). Figure 1. Scheme of the simulation model. The explosion occurs at the centre, and there is a density discontinuity at the plane . The LoS is in the xz plane with an angle of to the z -axis. After the supernova explosion, the spherical ejecta with a radius of R ej , a mass of M ej , and a kinetic energy of E ej is set at the centre of the simulation. Initially, the material in the ejecta has a velocity of , where v 0 is the velocity at the boundary of the ejecta (Jun Norman 1996 ; Truelove McKee 1999 ). The inner part of the ejecta within a radius of r c is homogenous with a density of c , whereas the outer part has a mass of M ej and a density with a power-law dependence on r with an index of s (Jun Norman 1996 ), i.e. where 0 is the density at r R ej . s 7 and 3/7 are used for the Type Ia SNR in this paper. The other details of the initial condition for the ejecta were shown in the equations(1)(3) in Jun Norman ( 1996 ). Without including the effects of the radiative cooling and the particle acceleration, the dynamical properties of the
Paper: 2018MNRAS.477.3164M (fulltext)
NGC 4449 must have been ongoing for a significant fraction of the burst lifetime. Depending on the 3D distribution of the extended H i , this material may not have been expelled into the IGM, but may instead reside at extended radii around the galaxy and/or recycled back into the galaxy. For NGC 5253, the detected outflow surrounds much of the central star-forming region and is strongest in the south. Marlowe et al. ( 1995 ) measure velocities of the Hα detected shells of 35 km s −1 with physical extents of ∼1 kpc. Summers et al. ( 2004 ) find the extent of the X-ray emission to be somewhat smaller (∼0.5 kpc). However, using combined Chandra observations, we find a more extended region in the south of the galaxy that reaches ∼2.4 kpc in projection that cannot be explained by dynamical arguments and short star formation time-scales. These inconsistencies are alleviated when the longer star formation time-scales from the SFHs are considered. Furthermore, the SFR in NGC 5253 has declined by nearly a factor of 6 over the duration of the burst. The declining SFRs and small-scale outflow in NGC 5253 indicate that we are seeing the end of a wind phase (see also Fig. 7 ), rather than just the beginning of an outflow as has been previously suggested (Summers et al. 2004 ). For NGC 1569, the dynamical age of the outflow was previously estimated to be 10 Myr (Heckman et al. 1995 ; Martin 1998 ). Using this time-scale, Martin et al. ( 2002 ) estimate the amount of oxygen that could be in the wind of NGC 1569 and compare this to the amount of metals produced by star formation over a 10–20 Myr time-scale assuming a nucleosynthesis yield. Based on their best-fitting model of a solar metallicity in the wind, and to maintain consistency with the current gas-phase oxygen abundance in the galaxy, these authors report nearly all metals formed in the recent star formation activity have been expelled in the wind. However, as stated in Martin et al. ( 2002 ), this requires that the present-day gas-phase oxygen content in the galaxy was created in previous episodes of star formation that must have been quiescent (i.e. not energetic enough to drive an outflow). Based on the CMDs and the derived SFHs, there has been significant star formation in the last few 100 Myr in NGC 1569, with the peak star formation occurring 65 Myr. This suggests that the wind has been ongoing. Furthermore, the best-fitting model of a solar oxygen abundance in the hot wind does not take into account the uncertainties; a nearly equal χ 2 is found for 0.25 Z ⊙ , as well as for supersolar abundances for NGC 1569 (Martin et al. 2002 , see their table 5). The amount of metals ejected by the wind in NGC 1569 remains an open question. 6 CONCLUSIONS We have examined the diffuse soft X-ray emission from Chandra observations in six nearby starburst dwarf galaxies from the STARBIRDS sample (McQuinn et al. 2015a ) and compared the X-ray properties with the neutral gas distributions, SFHs quantified from
Paper: 2020MNRAS.493..603Z (fulltext)
since once the magnetic field twists into an overwhelmingly toroidal configuration, as measured in the comoving jet frame, it is very difficult for it to share additional energy with the particle flow (Begelman Li 1994 ). The above results apply to a pair-dominated jet; the effect of an ion component on radiatively cooled turbulence is outside the scope of this paper. The ions would effectively not cool at all, but would probably receive the majority of dissipated energy (Zhdankin et al. 2019 ), suggesting that the true value of σ could be much lower due to a hot ion component. A self-consistency condition that needs to be met in order for external IC to dominate over synchrotron cooling, , could be violated if δ ph or Γ j is too small. A cooled turbulence model is also attractive for explaining flares from BL Lac-type blazars, which lack an intense ambient radiation field. Here it is plausible that the jet emission is self-contained and likely due to the SSC process, and even larger values of σ (at least the σ associated with the electron or pair content) would be required. 6 CONCLUSION In this work, we conducted and studied the results from a series of PIC simulations of driven turbulence in relativistic collision-less pair plasma with external IC radiation. We focused on understanding the statistical properties of turbulent fluctuations and particle distributions during the statistical steady state, where injected energy from external driving is balanced by radiative energy losses. We performed a parameter scan over the steady-state magnetization 〈σ〉 and system size L /2〈ρ e 〉, to the extent that computational resources allowed. We now summarize our main results: Regardless of initial conditions, turbulence eventually acquires a rigorous statistical steady state (by heating or cooling) with an equilibrium temperature that is close to analytic predictions (equation 7 ) and can be controlled by adjusting the energy density of the external photon bath, U ph . Radiative turbulence thus provides an ideal scenario for studying the fundamental properties of kinetic turbulence. The properties of turbulence change from low 〈σ〉 to high 〈σ〉. The former looks like standard non-relativistic MHD turbulence, apart from the tendency to have rotational magnetic fluctuations. The latter has fluctuations that are highly relativistic, compressible, and electromagnetically dominated. The steady-state particle energy distributions are always close to the thermal (Maxwell–Jüttner) distribution, with only a modest non-thermal population (broadening the tail by a factor of ∼2 in energy). The non-thermal population is more temporally variable and energetically significant at high 〈σ〉. These steady-state particle energy distributions are well fit by analytic solutions to the Fokker–Planck equation describing stochastic particle acceleration, for simple forms of the momentum advection and diffusion coefficients (linear and quadratic in energy, respectively). The
Paper: 2017ApJ...842...94T (fulltext)
by these telescopes (e.g., Serabyn Lacy 1985 ; Serabyn et al. 1988 ; Lacy et al. 1991b ; Roberts et al. 1996 ; Zhao et al. 2009 , 2010 ). The stretched appearance and kinematics of the GCMS suggest the models in which the streamers are tentative structures with Keplerian orbits around Sgr A* (e.g., Serabyn et al. 1988 ; Zhao et al. 2009 , 2010 ). There are still alternate models to explain these properties (e.g., one-armed spiral; Lacy et al. 1991b ; Irons et al. 2012 ). The tidal force of Sgr must have a serious effect on the interstellar medium (ISM) in the vicinity of Sgr A*, i.e., the GCMS (e.g., Lacy et al. 1982 ). Furthermore, the strong Lyman continuum radiation from the OB and WR stars in the Central cluster ionizes the ISM rapidly in the region (e.g., Genzel et al. 1996 ; Paumard et al. 2006 ). A recent dust observation of the GCMS with ALMA suggested that there is a possible scenario for the formation of the Central cluster: a star-forming molecular cloud is falling from a region somewhat far from Sgr A* and supplies young stars and the ISM to the vicinity of Sgr A* (e.g., Tsuboi et al. 2016 ). However, it is not clearly demonstrated how the gas is fed to Sgr A* because the innermost part of the ionized gas has very weak intensity and complicated structures. We present new observational results on the ionized gas streams in the vicinity of Sgr A* based on the observation of the H42 α recombination line using the Atacama Large Millimeter/submillimeter Array (ALMA). The distance of the Galactic center is assumed to be 8 kpc in this paper. 2. Observation and Data Reduction We have observed the GCMS in the H42 α recombination line (85.6884 GHz) as a part of the ALMA Cy.1 observation (2012.1.00080.S; PI: M.Tsuboi), which is an ionized gas tracer. The entire ALMA observation consists of a 137-pointing mosaic of the 12 m array and a 52-pointing mosaic of the 7 m array (ACA), covering a area including the “Galactic center 50 km s −1 molecular cloud” and the GCMS in CS (97.980953 GHz), SiO (86.846995 GHz), H 13 CO (86.754288 GHz), and H42 α emission lines. The molecular cloud encompasses the most conspicuous star-forming region in the vicinity of Sgr A*. We will present the detailed description of the analysis and the full results in another paper. We have detected the recombination line in the GCMS and the Sgr A East H ii region, and we concentrate on the GCMS in the H42 α recombination line in this paper. The data of the recombination line have angular resolutions of and using “natural weighting” and “Briggs weighting ( R = 0.5)” as u–v sampling, respectively. The frequency channel width is 244 kHz. The velocity resolution is 1.7 km s −1 (488 kHz). J0006−0623, J1517−2422, J717−3342, J1733−1304, J1743−3058, J1744−3116, and J2148+0657 were used as phase calibrators. The flux density scale was determined using Titan, Neptune, and Mars. Because the observation has a large time span of 19 months, the absolute flux density accuracy may be as
Paper: 2019MNRAS.482.5336S (fulltext)
of the crustal currents and the polar magnetic field. We compute the consequent X-ray luminosities and compare them with the observed values for magnetars. In Section 3 , we examine the dependence of conductivity tensor and Ohmic dissipation time-scale for high magnetic fields. In Sections 4 and 5 , we look at the X-ray and dipolar radiation luminosities, periods, and decay time-scales in the dynamo model and the screened core model. We find that in the screened core model, the ratio of the dipole radiation loss to the X-ray luminosity ( L X ) is a ‘U’ curve as a function of increasing spin-down age and period. Finally, in Section 6 , we try and understand the observed cut-off in the period of magnetars and end with a discussion in Section 7 . 2 CRUSTAL CURRENT AND POLAR MAGNETIC FIELD DISSIPATION AND X-RAY LUMINOSITY FOR SOME MODELS 2.1 Estimation of crustal magnetic field energy decay rate from the magnetar spin-down age Let us first make a simple estimate of the rate of magnetic energy dissipation in the crust. We assume the crust to be a shell of thickness, Δ R cr , in which the magnetic fields diminish with time as the crustal currents dissipate. We shall give some estimates which, though qualitative, capture some of the physics. The amount of magnetic energy stored in the crust is given roughly by the expression where B is the observed surface polar field and V is the volume of the crust. V = 4 R 2 · Δ R cr ∼ 4 × 10 17 cm 3 , where we have taken R ∼ 10 6 cm as the radius of the star and Δ R cr ∼ 1 km as the thickness of the crust. We shall use this definition of the crust for the rest of the paper. As the field drops to a fraction of its original value, we assume this energy will be dissipated and released as X-ray luminosity. We further assume that the typical time-scale in which this happens is given by the spin-down age of the magnetars, Thus, a crude estimate for the ‘average’ X-ray luminosity is given by dividing the stored magnetic energy by the spin-down age, This equation can be found in Dass Soni ( 2012 ). Hence, For , we use the relation below which is based on the dipole formula, As a standard (average) value of the polar field for magnetars we take, B p ∼ 4 × 10 14 G (which is the value of the initial field in the following Section 2.2 ). For such fields, the magnetar data indicates average periods of ∼8 s. On substitution, this yields an average crustal X-ray luminosity: The instantaneous luminosity when the periods reach ∼8 s will, of course, be less than the average. The dipolar spin-down radiation energy loss rate for magnetars with such magnetic fields and periods is then given by: where we have taken, I = 10 45 gm/cc, here and in what follows. 2.2 Estimation of magnetic field energy decay rate due to Ohmic resistance and the Hall effect The work of Gourgouliatos Cumming ( 2014 ) and Wood Hollerbach ( 2015 ) computes the magnetic field evolution in the crust via the Hall effect. In their work, they start with polar fields
Paper: 2016MNRAS.459.2603B (fulltext)
use the results of a double-Sérsic fitting to the surface density profile of our simulated galaxy, as described in Section 4.2 (the values are given in Table 2 ). Eris BH sits well below both relations, even though the black hole seeds start off close to the extrapolation at lower masses of the MBH-stellar mass relation (coloured bullets in the left-hand panel): this is because the black hole accretion rates are much lower than what would be required to keep up with the star formation rate in the central region (see Fig. 3 ). Figure 7. M BH – M Star and M BH – M Bulge relations. The red and blue symbols are from a compilation of Erwin Gadotti ( 2012 ) for early-type and late-type galaxies, respectively (with fits for the M BH – M Bulge relation shown by the dashed lines). The orange diamond is the location of the Milky Way in the plane as reported by Erwin Gadotti ( 2012 ). The yellow circle indicates the location of Eris BH . The coloured bulletts in the M BH − M Star figure show the location of the seeds, colour-coded as in Fig. 1 Table 2. Summary of global properties of Eris BH and Eris at z = 0. The stellar and gas masses in columns 1 and 2 refer to the total stellar and gas masses within the halo, also used to calculate the baryon fractions given in column 3. The disc and bulge masses have been obtained fitting a double Sérsic profile to the face-on surface density profile, with scale lengths given in columns 8 and 9. The last column gives the Sérsic index for the bulge, while the Sérsic index for the disc has been fixed to unity. Masses are given in solar masses, distances in kpc and velocities in km s − 1 . Having only one simulation, we can not give quantitative predictions on whether black holes in pseudobulges follow a different relation with the properties of the host (or if they do not follow any relation at all). But we note that our results on the M BH –σ and M BH – M Bulge relations are both consistent with what is found by Greene et al. ( 2008 ), and could also fit within the picture drawn by Kormendy et al. ( 2011 ), given the slow Seyfert-like growth of our simulated black hole. 4 BLACK HOLE FEEDBACK AND ITS EFFECTS ON THE HOST GALAXY As described in Section 2.2.2 , the simulation includes also a prescription for black hole feedback: a fraction of the energy extracted in the accretion process couples thermally with the surrounding medium. This standard prescription has been used extensively in other works, and its effects led to the idea of a ‘self-regulation’ of black hole growth to explain the origin of the M BH –σ relation and the ejection of gas from the host galaxy and the resulting quenching of star formation (e.g. Di Matteo et al. 2005 ). Given that the accretion rates are generally quite small for the central black hole of our simulated galaxy, we expect the effects of feedback to be modest. We find, in fact, that the amount of energy released by the black hole is significantly lower than the one released by
Paper: 2019A&A...629A..78C (fulltext)
2003 ). Its depletion covers a relatively narrow range of values ( Fig. 1 ) showing that it is not a strong function of environmental density. It has been hypothesized that the majority of carbon should be locked in graphite grains, as a likely explanation for the 2175 Å emission feature ( Draine 1989 , 2003 , and references therein). Graphite has been commonly adopted in ID models (e.g., Mathis et al. 1977 ). However, observational evidence pointed out that graphite cannot explain the variability of the 2175 Å feature ( Fitzpatrick Massa 2007 ). Furthermore, in analogy with the silicates, which are found to be amorphous, also graphite was deemed unlikely to survive in large quantities in the harsh environment of the ISM. Graphite should therefore face a natural process of amorphization ( Compiègne et al. 2011 ). The idea of carbon as a single and separate phase from the silicate population does not agree with a scenario of a constantly evolving and mixing medium ( Jones et al. 2017 ). Hydrogenated amorphous carbon (HAC) may indeed coat the silicate grains, forming a single population (e.g., Duley et al. 1989 ) with different characteristics with respect to the environment in which they reside and depending on the particle size (e.g., Jones et al. 2017 , and references therein). Polarization studies however have so far not confirmed this scenario. The carbon feature at 3.4 μ m shows a negligible degree of polarization with respect to the silicate feature at 9.7 μ m, pointing to two distinct grain populations ( Whittet 2011 ). Finally, under special condition of high pressure, for instance in a shocked environment, graphite and amorphous carbon can turn into nano-diamonds, which can constitute as much as 5% of the amount of C in the ISM ( Tielens et al. 1987 ), possibly with H and N inclusion ( Van Kerckhoven et al. 2002 ; Bilalbegović et al. 2018 ). Diamonds of possible ISM origin have been found in meteorites ( Lewis et al. 1987 ). An important carbon carrier are polycyclic aromatic hydrocarbons (PAH), large molecules (Ångstrom-sized) formed by carbon and hydrogen in a honeycomb structure. These molecules constitute up to about 10% of the carbon abundance (e.g., Tielens 2013 ). PAHs are very sensitive to ionizing radiation from far-ultraviolet to X-rays and they are easily destroyed near star formation sites at AU distance scale (e.g., Siebenmorgen Krügel 2010 ), up to kiloparsec scale, for active galaxies ( Voit 1992 ). Apart from C, other constituents can be studied in detail by future generation telescopes. Sulfur in the dust phase seems to be absent from the diffuse ISM ( Sembach Savage 1996 ). However, a relative fast transition to a depletion approaching −1 dex is reported in dense media, such as molecular clouds ( Joseph et al. 1986 ). In molecular clouds, sulfur can be included in aggregates such as H 2 S, SO 2 , OCS, SO, H 2 CS, NS SiS, CS, HNCS CH 3 SH ( Duley et al. 1980 , and references therein) as well as other
Paper: 2017MNRAS.464..754R (fulltext)
), the inner radius of the disc is 0.06 au. Additionally, the disc is in hydrostatic equilibrium in the vertical direction and its density profile ρ( r , θ) is given by the equation from Dullemond ( 2000 ): Here, Σ 0 is the superficial density at the outer edge of the disc, R 0 is the external radius of the disc and H ( r ) is the disc scaleheight. The last variable is defined from the self-irradiated passive disc proposed by Chiang Goldreich ( 1997 ) and is given by H ( r ) = r H 0 / R 0 ( r / R 0 ) 2/7 . The ratio H 0 / R 0 0.1 defines cold discs whereas H 0 / R 0 > 0.1 defines hot discs according to Faure, Fromang Latter ( 2014 ). Additionally, typical values for protoplanetary discs range from 0.1 to 0.3 (Ono, Nomura Takeuchi 2014 ). Hence, for a generic case, we adopted the value of 0.17. 4.2 Dust and ice model Because the interaction between light and matter in astrophysical environments is regulated by absorption and scattering processes, we calculated the opacity for a core of silicate 1 (Weingartner Draine 2001 ) with the inclusion of amorphous carbon and ice (H 2 O:NH 3 :CO 2 :CH 4 ) to simulate the presence of a mantle. The carbon inclusion was assumed to be 0.15 (Pontoppidan et al. 2005 ), and the abundances of ice relative to H 2 were 10 −4 for water and 10 −5 for CO 2 , CH 4 and NH 3 . To perform the mixing rule, we adopted the Maxwell–Garnett effective medium theory (Bohren Huffman 1983 ) to calculate the effective CRI for the grains. Considering that the relation between the complex dielectric function () and CRI () for non-magnetic materials is given by , the equation for the effective CRI () is given by where and are the CRI for the silicates and for inclusions, which can be carbonaceous material or ices, respectively. The opacity for the grains was calculated employing the Mie theory (Mie 1908 ), assuming spherical grains with a single grain size of 0.3 μm. Figs 6(a) –(c) show the opacities for the grains considering the Fluence 0, Fluence I (10 11 ions cm −2 ) and Fluence II (10 12 ions cm −2 ). As can be seen in Fig. 6 , the CRI presented in this paper allows us to calculate an opacity data base to be used in astrophysical models. This type of data base has been created by Ossenkopf Henning ( 1994 ), although such data are consistent only for non-processed regions in the ISM. However, the processing of the ice is a common scenario and the opacities from Ossenkopf Henning ( 1994 ) lead to wrong conclusions about the chemical evolution of the ice. In this way, the data base presented in this paper can be used to calculate new opacity tables, by using realistic data of processed ices, as well as realistic grain geometry. Figure 6. Absorption and scattering opacities of astrophysical grains employed in this paper. Panel (a) presents the opacities of bare silicate and panels (b) and (c) present the opacities of silicate with ice at Fluence 0, I and II. The lines were moved arbitrarily to highlight differences due to energetic