adsabs/FOCAL huggingface dataset:
It has been also suggested that the X-ray emission from SNe is dominated by inverse Compton (IC) scattering of photospheric optical photons by relativistic electrons accelerated by the SN shock on a timescale of weeks to a month after the explosion (see Chevalier & Fransson 2006). The X-ray luminosity strongly depends on the structure of the SN ejecta (ρSN ∝ r− n), the density structure of the CSM and the relativistic electron distribution responsible for the up-scattering (Chevalier & Fransson 2006; Margutti et al. 2012). Assuming a wind-like CSM (ρCSM ∝ r-2) and an ISM-like CSM (ρCSM = constant), Margutti et al. (2012) developed two generalized formalisms to calculate the IC luminosity. Their formalisms are strongly sensitive to the SN bolometric luminosity, Lbol (see Appendix A of Margutti et al. 2012). Here, it is assumed electrons are accelerated according to a power-law distribution n(γ) ∝ γ− p with index p = 3, as suggested by the radio observations of SN shocks in Type Ib/Ic explosions (see Soderberg et al. 2006; Margutti et al. 2014). Then, an IC luminosity in the wind-like CSM case is (see Margutti et al. 2012, their Appendix A) (1)\begin{eqnarray} \label{eq:1}  \frac{\mathrm{d} L_{\rm{IC}}}{\mathrm{d} \nu} \approx 2.1\times10^{-2}\left(\frac{\epsilon_{\rm e}}{0.1}\right)^{-2}\left(\frac{M_{\rm{ej}}}{1.4\,M_{\odot}}\right)^{-0.93}\left(\frac{A}{\rm{g\,cm^{-1}}}\right)^{0.64}  \nonumber \\\times\left(\frac{E}{10^{51}\,\rm{erg}}\right)^{1.29}\left(\frac{t}{\rm{s}}\right)^{-1.36}\left(\frac{L_{\rm{bol}}}{\rm{erg\,s^{-1}}}\right)\,\nu^{-1}, \end{eqnarray}dLICdν≈2.1×10-2ϵe0.1-2Mej1.4 M⊙-0.93Ag cm-10.64×E1051 erg1.29ts-1.36Lbolerg s-1 ν-1,where A = Ṁ/ (4πυwind), and ϵe is the fraction of thermal energy in the shock used for the electron acceleration (~0.1 see Chevalier & Fransson 2006). Theoretical IC X-ray luminosities with different mass-loss rates (for a wind velocity of 100 km s-1) are compared to the observations in Fig. 5. Here, both the observed bolometric light curve of SN 2005hk and the predicted bolometric light curve for the “N5def model” (see Fig. 8 of Kromer et al. 2013) are used as input for Lbol. The “N5def model” was shown to provide a good fit to the observations of the SN 2005hk event (Kromer et al. 2013). Moreover, E = 1.34 × 1050 erg, and Mej = 0.37 M⊙ (these parameters are consistent with the N5def model) are used in Eq. (1) to estimate the X-ray luminosities. Here, we note that SNe Iax have a range of luminosities, ejecta velocities, and inferred ejecta masses (Foley et al. 2013). Different bolometric luminosity light curves and ejecta masses of different SNe Iax may change the results shown in Fig. 5.
Citation: Margutti et al. 2014
Citation function: Uses
Our model consists of four zones; an inner disk (zone 1, as in Matter et al. 2016), and three zones into which the outer disk is divided (zones 2–4) in order to produce the azimuthal asymmetries seen in our SPHERE observations. The radial extent of these zones were constrained from our images: zone 2 corresponds to the component inside ~14 au observed in the noncoronagraphic J-band image, zone 3 corresponds to the bright ring at 16 au, and zone 4 corresponds to the asymmetrical outer disk. The radial extent of each of these zones, as well as their disk masses, is summarized in Table 2. No gaps were introduced between the three outer disk zones, but the gap between zones 1 and 2 from the Matter et al. (2016) model was kept. We note that these radii are only loosely constrained from our data, and are determined by eye from both the coronographic and noncoronographic data. We ran a grid of models with different inclinations and position angles for zones 1–2 (note that we systematically use the same inclination and position angle for these two zones) and for zone 3 – that is, a total offour free parameters, keeping in mind that the relative inclination between components must be small to allow a single broad shadow to be cast (as opposed to two narrow shadow lanes). The inclination and position angle of zone 4 is set to the values provided in Sect. 3. The grid was sampled in steps of 1° for inclination, and 2° for PA, and later refined to 0.5° steps for inclination and 0.5° for PA after a good initial agreement is found between the average azimuthal profile of the model and the scattered light images. The best fitting model was picked not only based on the location of the shadows cast by the inner components on the outer disk, but also on the shape (slope) of the resulting azimuthal profile of the outer disk. There is also a degeneracy if we consider that the inclinations and PAs of the two inner components (zones 1–2 and zone 3) can be exchanged and produce very similar results. However, doing this would cast a shadow in a different location, and thus produce a different azimuthal profile for zone 3.
Citation: Matter et al. (2016)
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After the LVC announcement, a multi-wavelength campaign immediately started. The campaign involved the X-ray and gamma-ray satellites first, and radio/IR–optical/TeV ground observatories at later times. The detection of a new optical transient (OT) was first announced by the One-meter, Two-hemisphere (1M2H) team discovered with the 1 m Swope telescope on August 18 01:05 UT (Coulter et al. 2017a, 2017b; Drout et al. 2017), named Swope Supernova Survey 2017a (SSS17a, now with the IAU designation AT 2017gfo). The OT was also detected independently by five other teams, the Dark Energy Camera (Allam et al. 2017; Cowperthwaite et al. 2017), the Distance Less Than 40 Mpc Survey (L. Tartaglia et al. 2017, in preparation; Valenti et al. 2017; Yang et al. 2017), Las Cumbres Observatory (Arcavi et al. 2017a, 2017b, 2017c), the Visible and Infrared Survey Telescope for Astronomy (Tanvir et al. 2017a, 2017b), and MASTER (Lipunov et al. 2017b, 2017a), REM-ROS2 (Melandri et al. 2017; Pian et al. 2017a, 2017b), Swift UVOT/XRT (Evans et al. 2017b, 2017a), and Gemini-South (Kasliwal et al. 2017; Singer et al. 2017), and for all see also MMA17. AT 2017gfo is located at 106 from the early-type galaxy NGC 4993, at a distance of ∼40 Mpc. A sequence of satellite high-energy observations started almost immediately, with exposures of the GW170817 LR depending on satellite position and operations. The X-ray and γ-ray observations included contributions by CALET (Nakahira et al. 2017), Konus-Wind (Svinkin et al. 2017), Insight-HXMT (Li et al. 2017; Liao et al. 2017), AstroSat CZTI (Balasubramanian et al. 2017), AGILE-GRID (see below), Fermi-LAT (Kocevski et al. 2017), MAXI (Sugita et al. 2017; S. Sugita et al. 2017, in preparation), Super-AGILE (MMA17; this work), Swift X-ray Telescope (Evans et al. 2017b, 2017a), NuSTAR (Harrison et al. 2017), INTEGRAL JEM-X (Savchenko et al. 2017b; V. Savchenko et al. 2017, in preparation), and Chandra (Fong et al. 2017; Margutti et al. 2017a; Troja et al. 2017b). An X-ray counterpart detection at the OT position was reported after 9 days and confirmed after 15 days by Chandra (Fong et al. 2017; Margutti et al. 2017b; Troja et al. 2017a, 2017b). Moreover, an important detection in the radio band has been reported by VLA (Alexander et al. 2017a, 2017b).
Citation: Tanvir et al. 2017a
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Cassiopeia A (Cas A), which is a remnant of an SN IIb explosion (Krause et al. 2008), is an ideal target for examining the properties of SN dust. It is quite young (∼330 years; Fesen et al. 2006), so the SN ejecta and freshly formed SN dust have not been significantly mixed with the circumstellar/interstellar medium (CSM/ISM), and it is relatively close (3.4 kpc; Reed et al. 1995), so the physical and chemical structures can be resolved. Therefore, Cas A has been a major target of infrared space missions. The first direct evidence for SN dust in Cas A came from the Infrared Space Observatory (ISO), which performed mid-infrared (MIR; 2.4–45 μm) spectroscopic observations toward the bright ejecta shell (Lagage et al. 1996; Arendt et al. 1999; Douvion et al. 2001). The continuum spectra had a strong bump peaking at 21 μm together with a relatively weak bump at 9.5 μm, and Douvion et al. (2001) showed that the spectra can be fitted by two dust components at different temperatures, one at 90 K and the other at 350 K, with pyroxene (MgSiO3), quartz (SiO2), and aluminum oxide (Al2O3) as major components. The composition and distribution of this “warm” SN dust heated by a reverse shock have been studied in detail using the spectral mapping data of the Spitzer Space Telescope (Ennis et al. 2006; Rho et al. 2008; Arendt et al. 2014). These studies showed that the dust emission exhibits distinct spectral characteristics depending on the SN material with which it is associated. The most prominent dust emission is that with strong 9 and 21 μm bumps, which is associated with SN ejecta having strong Ar emission lines. The spectral features agree with the ISO spectra and can be reproduced by Mg protosilicate/MgSiO3, Mg





SiO





, or SiO2. The dust emission associated with SN ejecta having strong Ne lines is smooth without any silicate features and can be reproduced by Al2O3 dust or C glass. There is yet another type of smooth dust emission associated with X-ray Fe emission, but it is probably from swept-up CSM (Arendt et al. 2014). Meanwhile, Nozawa et al. (2010) showed that the observed spectral energy distribution (SED) of dust emission can be fitted by a physical model in which the SNR is expanding into a dense CSM, and the MIR emission is mostly from silicate grains (e.g., MgSiO3, Mg2SiO4, and SiO2) and MgO. The estimated total mass of warm dust ranges from 0.008 to 0.054 





.
Citation: Lagage et al. 1996
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These data were compared to two different grids of stellar evolution models (see Table 1 for a summary). The first was a grid run using the Yale Rotating Evolution Code (YREC, Pinsonneault et al. 1989, with updates as discussed in van Saders & Pinsonneault 2012). We ran a grid of masses (0.6 M⊙–2.6 M⊙ in 0.1 M⊙ increments), metallicities ([Fe/H] = −2.0 to +0.6 in steps of 0.2), and α-element enhancements (+0.0, +0.2, and +0.4) to generate tracks. These models use a Grevesse & Sauval (1998) elemental mixture, with the updated OPAL equation of state (Rogers et al. 1996; Rogers & Nayfonov 2002), OPAL opacity tables (Iglesias & Rogers 1996) and gray atmospheres. Alpha-enhanced models use alpha-enhanced starting models and opacity tables, but not equations of state. Our models include semiconvection (see Kippenhahn & Weigert 1994), but do not include overshoot or the diffusion of helium or heavy elements. We assumed a Big Bang helium value of Y = 0.2485 at Z = 0 (Cyburt 2004), and did a linear fit to our solar value of Y = 0.272683 at Z = 0.018011 (so 





). We convert Z to metallicity assuming 





 for a solar mixture, with a correction factor for alpha-enhanced models. We compare to other theoretical models in Section 3. In the initial part of this analysis, we use grids with a solar mixing length (Böhm-Vitense 1958) of 1.72, chosen to reproduce the solar luminosity (3.827 × 1033 erg s−1) and the solar radius (6.957 × 1010 cm) at the solar age (4.57 Gyrs) (Mamajek 2012). For later parts of this study, we interpolate in a grid of models with mixing lengths of 1.22, 1.72, and 2.22. For the analysis in Section 3.2, we varied the initial helium abundance and surface boundary conditions in our model grid to explore their impact on the theoretical uncertainties. We ran models with fixed helium fractions of 0.239, 0.290, and 0.330, in addition to the solar calibrated value of 0.272683. We also created models with different atmospheric surface boundary conditions (defined as the boundary pressure at τ = 2/3): a look-up table from Kurucz (1997) models, and one from Castelli & Kurucz (2004) models, in addition to the gray atmosphere used in our main model grid.
Citation: Iglesias & Rogers 1996
Citation function: Uses; Uses
The distances to planetary nebulae (PNe) have always faced the difficulty that nearby targets were lacking that could be reached well by direct methods. Trigonometric parallaxes have been obtained in a homogeneous long time-line campaign by the US Naval Observatory (USNO; Harris et al. 2007) and from the Hubble Space Telescope (HST; Benedict et al. 2009). Other studies (Acker et al. 1998; Smith 2015) showed that Hipparcos spacecraft parallaxes do not seem to be reliable. It was assumed that contamination by the emission of the surrounding nebulae caused these problems. Another model-independent method for distances to PNe are a cluster membership, as studied extensively by Majaess et al. (2007, 2014), and as discussed in Frew et al. (2016). In addition to these model-independent methods, a wide variety of statistical, model-dependent individual distance scales have been derived. The most frequently used of these are certainly those that are based on surface brightness versus angular sizes. They sometimes include optical depth corrections. All these methods have to be calibrated against a data set of nebulae with known distances. The older, widely used method is based on the 6 cm radio continuum flux, either using the ionized mass concept of Daub (1982) in the calibrations of Cahn et al. (1992) and Stanghellini et al. (2008), or by means of the radio continuum brightness temperature as used by van de Steene & Zijlstra (1994) and calibrated with a Galactic bulge sample. The newest model developed by Frew et al. (2016) is based on similar ideas, but makes use of the optical Hα surface brightness and a wide set of various calibrators. Moreover, they use a completely homogeneous data set for the brightness data derived earlier by themselves (Frew et al. 2013). Smith (2015) and Frew et al. (2016) described the underlying physics and assumptions for all these methods in detail. With the upcoming Gaia project (Gaia Collaboration 2016), a new era was expected to start for many classes of objects. The first step into this was described by Stanghellini et al. (2017) based on the combined Tycho + Gaia DR1 solution called TGAS (Michalik et al. 2015). With the second Data Release of Gaia (hereafter GDR2; Gaia Collaboration 2018), a complete homogeneous data set based only on Gaia measurements is available now for the first time. We present here the comparison of this new data set with common previous calibrations of PNe distances. Moreover, we compare it to the preliminary TGAS results in Stanghellini et al. (2017). Finally, we discuss possible caveats using the current GDR2.
Citation: Frew et al. (2016)
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Suprathermal electron strahls in the solar wind come from the Sun and are focused along magnetic field lines (Feldman et al. 1975; Rosenbauer et al. 1977; Pagel et al. 2005). Therefore, observations of CSE strahls within MCs can indicate that the flux rope structures are still connected with the Sun’s magnetic field lines on both ends (Gosling et al. 1995; Larson et al. 1997; Shodhan et al. 2000; Feng et al. 2015, 2019). Counterstreaming suprathermal electrons can also be produced by other mechanisms, for example connection to the Earth’s bowshock (Feldman et al. 1982; Stansberry et al. 1988), interplanetary shocks or corotating interaction regions (CIRs; Gosling et al. 1993; Steinberg et al. 2005; Lavraud et al. 2010), and distribution function depletions near the 90° pitch angle (Gosling et al. 2001, 2002; Skoug et al. 2006). Among these mechanisms, the depletion CSEs are often observed on closed or open field lines within ICMEs, but the depletion CSEs are centered on and roughly symmetric to the 90° pitch angle (Gosling et al. 2002), and can be distinguished from CSE strahls. Shodhan et al. (2000) examined the CSE strahl signatures of 52 MCs detected by a spacecraft near 1 AU and determined that approximately 87.5% of MCs exhibit CSE signatures, revealing that most MCs are still attached to the Sun at both ends at 1 AU. Gosling (1990) and Gosling et al. (1995) proposed explanations for how flux ropes arise in terms of three-dimensional reconnection close to the Sun. They illustrate how the original flux rope reconnects to form a helical field line connected to the Sun at both ends. The closed field lines of flux ropes can gradually open and occasionally disconnect from the corona when closed flux ropes expand from the Sun into the interplanetary space (Gosling et al. 1995). The proposal of Gosling et al. (1995) was confirmed by statistical results of Shodhan et al. (2000), namely most MCs exhibit CSE strahls in parts of their durations, and only 6 of the 52 MCs have no CSE. We therefore want to know whether or not nonMC ICMEs have the same CSE strahl signatures. In this study, CSE signatures from the Advanced Composition Explorer (ACE) from 1998 to 2008 are compared between ICMEs with and without MCs, and we discuss whether the CSE signatures are related to the flux-rope structures.
Citation: Lavraud et al. 2010
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We used the theoretical models by Eggleton et al. (2008), Charbonnel & Lagarde (2010), and Lagarde et al. (2012) for the comparison. They provide quantitative values representing the first dredge-up, thermohaline (TH), and thermohaline and rotation (TH+V) induced mixing. Eggleton et al. (2008) estimated the mixing speed with their formula for the diffusion coefficient and found that a range of three orders of magnitude in their free parameter can lead to the observed levels of 12C /13C. Their predicted value of the 12C /13C ratio for a solar-metallicity 2 M⊙ star at the RGB tip is 17. A more recent model of the thermohaline-induced mixing by Charbonnel & Lagarde (2010) lists for the same stars a higher value of about 20. The model of Charbonnel & Lagarde (2010) of thermohaline instability induced mixing and the model by Eggleton et al. (2008) are both based on the ideas of Eggleton et al. (2006) and Ulrich (1972). It includes developments by Charbonnel & Zahn (2007). Eggleton et al. (2006) found a mean molecular weight (μ) inversion in their 1 M⊙ stellar evolution model, which occurred after the so-called luminosity bump on the RGB, when the hydrogen-burning shell reaches the chemically homogeneous part of the envelope. The μ-inversion is produced by the reaction \hbox{$^3{\rm He(}^3{\rm He}, 2p)^4{\rm He}$}He(33He,2p)4He, as predicted in Ulrich (1972). It does not occur earlier because the magnitude of the μ-inversion is low and negligible compared to a stabilising μ-stratification. Charbonnel & Zahn (2007) computed stellar evolution models including the ideas of Kippenhahn et al. (1980), who extended Ulrich’s equations to the case of a non-perfect gas. Charbonnel & Zahn (2007) also introduced a double diffusive instability (i.e. thermohaline convection) and showed its importance in the chemical evolution of red giants. This mixing connects the convective envelope with the external wing of the hydrogen-burning shell and induces surface abundance modifications in evolved stars (Charbonnel & Lagarde 2010).
Citation: Eggleton et al. (2006)
Citation function: Compare/Contrast
Here, we propose a new speculative site for strong r-process nucleosynthesis in which it is formed by jets launched by an NS spiralling-in inside the core of a giant star. These jets will also explode the star. This is a rare evolutionary route and hence complies with the finding of large variations in the abundances of these elements. In this first study, we limit ourselves to present the scenario and show its viability. Most ingredients of our newly proposed scenario were studied in the past but were never put together into a coherent picture to yield a new possible site for the strong r-process. Previously studied ingredients of our proposed scenario include the CE of an NS and a giant (Thorne & Zytkow 1975; Armitage & Livio 2000; Chevalier 2012), the launching of neutron-rich gas from accretion discs around compact objects (e.g. Surman & McLaughlin 2004; Kohri et al. 2005), the launching of jets by NS accreting at a high rate (Fryer, Benz & Herant 1996) and the formation of r-process elements in jets from NS (Fryer et al. 2006). The idea that jets can explode stars under specific conditions was also studied in the past, e.g. in a CE evolution (Chevalier 2012). Another specific class of models are based on magnetic amplification by a rapidly rotating stellar core (e.g. LeBlanc & Wilson 1970; Bisnovatyi-Kogan, Popov & Samokhin 1976; Meier et al. 1976; Khokhlov et al. 1999; MacFadyen, Woosley & Heger 2001; Woosley & Janka 2005; Couch, Wheeler & Milosavljević 2009). This magnetorotational mechanism creates bipolar outflows (jets) around the newly born NS that are able to explode the star. However, the required core's rotation rate is much larger than what stellar evolution models give, hence making most of these models applicable for only special cases. We, on the other hand, claim that all CCSNe are exploded by jets, the jittering-jets model, that also synthesize r-process elements (but not the strong r-process), and hence our newly proposed scenario is part of a unified picture we try to construct for exploding all massive stars and the synthesis of r-process elements. In the jittering-jets model, the explosion of CCSNe is powered by jittering jets launched by an intermittent accretion disc around the newly born NS (Papish & Soker 2014a,b). The intermittent accretion disc is formed by gas accreted from the convective core regions that have a stochastic angular momentum (Gilkis & Soker 2014; Papish & Soker 2014a).
Citation: Kohri et al. 2005
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The dynamic viscosity of a fully ionized gas in the absence of magnetic fields (see Spitzer 1962)1(2)
\begin{equation}
\mu _{\rm Spitzer} = 2.21\times 10^{-15}\,\frac{(T\,[\mathrm{K}])^{5/2}A_{\rm i}^{1/2}}{Z^4\ln \Lambda }\,\mathrm{g\,cm^{-1}\,s^{-1}}\,
\end{equation}
is mainly a function of the temperature T. The parameters Ai and Z are the atomic weight and charge, respectively, of the ions. The Coulomb integral ln Λ depends logarithmically on temperature and density, and is commonly approximated by a constant value around 40 for clusters (see Gaspari & Churazov 2013; Smith et al. 2013). Slices of the resulting viscosity are shown in middle panels of Fig. 1 for three stages of the simulated merger. In Fig. 2, radial profiles of volume averages (blue), medians (red) and interquartiles (red shaded) are plotted for the cluster produced by the merger at z = 0. The typical value of the Spitzer viscosity in the WHIM is μSpitzer ∼ 10 g cm−1 s−1, with a gradual increase towards the core. Around the accretion shock (R ≈ 5 Mpc), the viscosity drops rapidly. It turns out that the median profile reflects this drop much better than the volume-averaged (or mass-averaged) profile because averaging a quantity that varies over several orders of magnitude is biased towards high values. Beyond the outer shocks, the values calculated with equation (2) become physically meaningless because the unshocked gas is neutral or only partially ionized. For the ICM, viscosities above 100 g cm−1 s−1 would be representative if it were not for the depression within 100 kpc from the centre. As discussed in Schmidt et al. (2016), this is a consequence of the overpronounced cool cores in our simulations. Feedback, particularly from AGNs, would raise the core temperature and, thus, the viscosity. This is important to bear in mind for the following discussion: In any case, the viscosity following from equation (2) is large enough to significantly affect or even suppress dynamical instabilities in both the ICM and the WHIM. This was also demonstrated by Roediger et al. (2013), Roediger et al. (2015) and ZuHone et al. (2015) in the context of Kelvin–Helmholtz instabilities and gas sloshing induced by minor mergers.
Citation: Gaspari & Churazov 2013
Citation function: Uses