adsabs/FOCAL huggingface dataset:
Recently, another explosion mechanism has gained more attention in the context of both the explosion dynamics and formation process of neutron stars composed of exotic matter at their interiors (see Fischer et al. 2018; Zha et al. 2020). When the energy density becomes sufficiently high, then nuclear—in general, hadronic—matter undergoes a phase transition to the deconfined quark–gluon plasma. Therefore, the solution of quantum chromodynamics (QCD)—the theory for strong interactions with quarks and gluons as the degrees of freedom—predicts a smooth crossover transition at a pseudocritical temperature in the range of 150–160 MeV (Bazavov et al. 2014; Borsányi et al. 2014; Bazavov et al. 2019). However, this is the case only at vanishing baryon density, and hence cannot be applied to astrophysical applications of compact stellar objects such as neutron stars and supernovae. At high density, phenomenological quark matter models have long been employed in astrophysical studies, based on the common two-phase approach with separate hadronic and quark matter phases. Such phase transition construction results not only in the transition to the third family of compact stars, known as hybrid stars, as an extension to the neutron stars second family, but also in the as-yet incompletely understood question of the nature of the QCD phase transition at high baryon density (see Annala et al. 2020, and references therein). It is likely that low- and intermediate-mass NSs, having masses of ∼1.4 M
⊙, can be explained solely by hadronic equations of state (EOS). More massive NSs with ∼2 M
⊙, the presence of which has been observationally confirmed at high accuracy (see. Demorest et al. 2010; Antoniadis et al. 2013; Cromartie et al. 2020), feature the highest baryon densities encountered in the universe. It is a valid question to ask if a QCD transition occurs at the densities encountered at the interior of such massive pulsars (see Kurkela et al. 2014, and references therein). Furthermore, if confirmed, this is likely to affect our understanding of the core-collapse supernova phenomenology.
Citation: Bazavov et al. 2019
Citation function: Compare/Contrast; Compare/Contrast
The magnetic field configuration at the deep photosphere shows a low twist in most of the umbra (Fig. 4b). Only some penumbral regions exhibit a certain twist, which is generally below 30°. On the contrary, the twist measured at the height of the Si I 10 827 Å line shows high values at the bottom right quarter, which reach values up to −90° around the umbra-penumbra boundary. This twist contrasts with that found inside the umbra at the bottom left quarter, which is mainly positive. Socas-Navarro (2005) found that the magnetic twist can change the sign from the photosphere to the chromosphere. Our observations indicate that the twist can change from the deep photosphere (where it is almost negligible) to higher photospheric layers (where negative and positive twists coexist). This variation takes place over a height difference of around 130 km (Felipe et al. 2018). However, the degree of the twist may be sensitive to the choice of the center of the sunspot. In this context, it can be more meaningful to discuss the magnetic field torsion. Figure 4e shows that the bottom part of the penumbra has a slightly negative torsion, with stronger values near the umbra-penumbra boundary (in the same region where the twist from the Si I 10 827 Å line was higher). Since the magnetic field at the formation height of the Ca I 10 839 Å line is approximately pointing towards the center, a negative torsion indicates that at the formation height of the Si I 10 827 Å line the horizontal component of the magnetic field is shifted clockwise from the radial direction (see Fig. 5 for a vectorial representation of the horizontal magnetic field). A positive torsion, such as that measured at the left and bottom part of the umbra, means that the magnetic field is shifted counterclockwise. The position on the horizontal plane of the wavefront of a slow magnetoacoustic wave propagating along magnetic field lines from the deep photosphere to the high photosphere will rotate according to the direction of the torsion. As discussed in Sect. 3.1, the bottom arm of the spiral wavefront initially appears at the left side of the umbra and moves counterclockwise towards the center of the sunspot. This movement is in agreement with the positive torsion measured at that region from the analysis of the spectropolarimetric map. On the contrary, at the regions where the measured torsion is negative, we do not detect clockwise motions of the spiral wavefront. In that part we only find a radially outward movement.
Citation: Felipe et al. 2018
Citation function: Uses
The mergers of double neutron star systems or the neutron star-black hole binaries generate strong gravitational wave (GW) radiation as well as short-duration gamma-ray bursts (SGRBs; including the so-called long-short GRBs that have a duration longer than 2 s but are unaccompanied by supernova emission down to very stringent limits; Eichler et al. 1989; Piran 2004; Della Valle et al. 2006; Fynbo et al. 2006; Berger 2014). Before 2017, it was widely believed that the GW/SGRB association rate is low because the SGRB outflows are highly collimated with a typical half-opening angle of ∼0.1 rad (Clark et al. 2015; Li et al. 2016). Surprisingly, on 2017 August 17, the gamma-ray monitor on board the Fermi γ-ray space telescope and INTEGRAL had successfully detected a weak-short GRB 170817A (Goldstein et al. 2017; Savchenko et al. 2017) that is spatially and temporally correlated with GW170817, the first neutron star merger event detected by Advanced LIGO/Virgo (Abbott et al. 2017). The GW/SGRB association has been formally established. However, considering the relatively small event distance (D ∼ 40 Mpc), the isotropic-equivalent gamma-ray radiation energy of GRB 170817A is just ∼3 × 1046 erg, which is at least 100 times dimmer than that of the typical SGRBs. An underluminous SGRB could either result from the breakout of the mildly relativistic shock from the leading edge of the merger-driven quasi-isotropic sub-relativistic ejecta (Kasliwal et al. 2017), or be the faint prompt emission of a highly structured relativistic ejecta viewed at a large polar angle (Jin et al. 2018). The puzzling fact that GRB 170817A and the long-duration GRB 980425 (at a distance of D ∼ 36 Mpc, and that has been suggested to be the shock breakout signal; Kulkarni et al. 1998), the two close events with remarkably different progenitors, have rather similar luminosity and spectral peak energy (Wang et al. 2017), may favor the shock breakout model. It is thus unclear whether or not GW170817-like mergers are indeed the sources of the bright SGRBs. The forward shock afterglow observations of GW170817/GRB 170817A are helpful in answering such a question. Though the “early” rising X-ray and radio afterglow emission could be reproduced by a cocoon-like mildly relativistic ejecta (Kasliwal et al. 2017), the late-time afterglow data modelings strongly favor the presence of an off-axis relativistic (structured) outflow component (D’Avanzo et al. 2018; Lamb et al. 2018; Mooley et al. 2018b; Yue et al. 2018). Particularly, the off-axis relativistic outflow component at a viewing angle of θv ∼ 0.35 rad has been convincingly identified/measured in the radio image (Mooley et al. 2018a; Ghirlanda et al. 2019). Nevertheless, a direct “observational” link between GW170817/GRB 170817A and bright SGRBs is still lacking. In this Letter, we carry out statistical studies of the SGRB afterglow data and aim to establish such a connection.
Citation: Abbott et al. 2017
Citation function: Motivation; Motivation
In the coupled ADAF-jet model, the accretion flow ADAF and the jet are connected by a defined parameter, $\eta \equiv \dot{M}_{\rm jet}/\dot{M}$, and $\dot{M}_{\rm jet}$ is input by assuming a value of, η, which is free parameter in the present model. The half-opening angle ϕ of the jet in the low/hard state of NS-LMXBs is uncertain. In this paper, we fix ϕ = 0.1 as assumed by several other authors for modelling the SED of the BH-LMXBs (e.g. Yuan et al. 2005; Zhang et al. 2010). Observationally, the bulk Lorentz factor of the jet in the low/hard state of X-ray binaries can be restricted in a relatively narrow range and the velocity of the jet is mildly relativistic, i.e. Γjet ≲ 2. More strictly, the bulk Lorentz factor is restricted to be as Γjet ≲ 1.67 (Gallo et al. 2003), and Γjet ≲ 1.2 (Fender 2006). In the internal shock model, the energy density of the internal shock increases with increasing Γjet, which finally will result in an increase of both the radio emission and X-ray emission (Yuan et al. 2005). However, since Γjet is restricted in a very narrow range by observations, we expect that a slight change of Γjet will result in a slight change of the jet emission. In this paper, we fix Γjet = 1.2 corresponding the bulk velocity of the jet 0.55c (Fender 2006). The value of ϵe and ϵB describing the fraction of the internal energy of the internal shock stored in the accelerated electrons and the magnetic field, and the index, pjet, describing the power-law distribution of the electrons in the jet after the acceleration by the shock are uncertain. Qiao & Liu (2015) tested the effect of ϵe and ϵB on the emergent spectrum of the jet in an observationally inferred range of 0.01  ϵe  0.1 and 0.01  ϵB  0.1. It was found that a change of ϵB in the range of 0.01–0.1, the emergent spectrum of the jet nearly does not change (see the right-hand panel of fig. 3 of Qiao & Liu 2015). A change of ϵe in the range of 0.01–0.1, the radio spectrum nearly does not change. However, the X-ray luminosity changes obviously by changing the value of ϵe from 0.01 to 0.1 (see the left-hand panel of fig. 3 of Qiao & Liu 2015). As shown in the left-hand panel of fig. 2 of Qiao & Liu (2015), the X-ray emission is completely dominated by the accretion flow (corona) in the luminous X-ray state, which is also true in this paper, i.e. the X-ray emission from the ADAF completely dominates the X-ray emission from the jet. In this paper, we fix ϵe = 0.04 and ϵB = 0.02, respectively, throughout the paper as (Qiao & Liu 2015). The value of the power-law index pjet of the electron distribution in the jet predicted by the shock acceleration is 2  pjet  3. By modelling the SEDs of three BH-LMXBs, the value of the power-law index pjet of the electron distribution is constrained to be 2.1 (Zhang et al. 2010). Meanwhile, a change of pjet in the range of 2  pjet  3 has very minor effect on the X-ray spectrum. In this paper, we fix the power-law index pjet = 2.1 throughout the paper.
Citation: Fender 2006
Citation function: Uses
Another important sink for CO2 is the loss of CO2 by weathering of the surface. We neglected this effect in the present study as in Tosi et al. (2017). Foley & Smye (2018) study the potential impact of weathering for stagnant-lid planets. They find that for large enough planetary carbon reservoirs and radiogenic heating, weathering and outgassing can balance each other for 1000–5000 Myr, thus enabling habitable surface conditions. On the one hand, for large carbon reservoirs they show that weathering is supply-limited, i.e. the supply of fresh surface material is likely too small such that CO2 would accumulate in the atmosphere, which leads to a Venus-like, uninhabitable hot climate. On the other hand, for low outgassing rates limited by the carbon reservoir, the planet would likely exist in a snowball state. For the snowball state limit, they use outgassing values determined by Haqq-Misra et al. (2016) and Kadoya & Tajika (2014) as lower limits, which are equal to 10 and 100% of present Earth’s outgassing, respectively; 10% of Earth’s CO2 outgassing flux, as adopted by Foley & Smye (2018), corresponds to ~2.6 ⋅ 1010 kg yr−1. In our models the rate of CO2 outgassing depends on the initial water concentration and on the assumed redox state. For 62 ppm water in the mantle and 

$f_{\mathrm{O}_2}$



f


O
2





 at the IW buffer, the outgassing rate during the phase of melt production peaks at about 2.5 ⋅109 kg yr−1, while for 

$f_{O_2}$



f


O
2





 at IW+1 it is one order of magnitude higher, and hence comparable with the lower limit mentioned above. Foley & Smye (2018) discuss weathering for Earth-like water reservoirs. In our scenarios, the weathering of the surface could be limited by the smaller surface water reservoir and lower precipitation rates. Therefore, although our outgassing rates are close to the lower limits adopted by Foley & Smye (2018) that might lead to snowball states, the overall impact of weathering on our findings is not easily predictable and will require further study. Furthermore, Foley & Smye (2018) consider metamorphic outgassing as an additional source of CO2, which could increase CO2 outgassing rates. Metamorphic outgassing refers to the release of CO2 associated with the decarbonation of buried carbonated crust. On Earth, the contribution of this mechanism to the planet’s long-termcarbon cycle is thought to be potentially important, but is still poorly understood (Evans 2011). Furthermore, the release of carbon due to metamorphic processes is generally associated with partial melting in subduction zones (e.g. Dasgupta & Hirschmann 2010) and deformation in active orogenic regions (e.g. Evans et al. 2008), two settings that are intimately related to plate tectonics and hence absent in a stagnant-lid planet. Whether metamorphic outgassing on such bodies can be as relevant as on a planet with plate tectonics is therefore unclear and deserves more investigation.
Citation: Haqq-Misra et al. (2016)
Citation function: Compare/Contrast; Compare/Contrast
The hydrodynamic equations (1–5) are solved in this paper using total variation diminishing (TVD) scheme, introduced and developed by Harten (1983). The scheme (or the modified version of it) is applicable to hydrodynamic problems and has been used extensively in relevant astrophysical applications (Ryu 1993; Ryu et al. 1995a; Ryu, Jones & Frank 1995b; Lee, Ryu & Chattopadhyay 2011; Chattopadhyay et al. 2012; Lee et al. 2016). TVD scheme is an Eulerian, second-order accurate, non-linear, finite difference scheme, which accurately captures shock. The temporal and spatial evolution of the conserved quantities ρ, ρvi, and ${\cal E}$ is computed using approximate Roe type Riemann solver to solve the differential equations, followed by application of a non-oscillatory first-order accurate scheme to the modified flux functions to achieve second-order accuracy (see Roe 1981; Harten 1983; Ryu 1993). Equations of motion (1–5) are similar to those solved in Chattopadhyay et al. (2012). In Chattopadhyay et al. (2012), the galactic outflow was powered by the radiation from the galactic disc, while being decelerated by the gravity of the galactic disc, the halo, and the bulge matter. In contrast, in this paper, the accretion disc outflow is powered by the radiative fluxes and the centrifugal force from the KD, and is decelerated by the radiative drag terms as well as the gravity of the central black hole. To solve the equations of motion (1–5), we considered the TVD scheme (see Chattopadhyay et al. 2012, for details) for the resolution 512 × 512. A schematic representation of the computational arrangement is presented in Fig. 2, which marks the ghost cells where the boundary conditions are implemented and also the computational domain. We employed continuous boundary condition at z = 0 boundary, and outflow boundary condition at the outer r and z boundaries (i.e. no inflow but continuous if v > 0). At r = 0, or the axis of symmetry, reflection boundary condition has been employed. The type of boundary conditions employed are also mentioned in Fig. 2. We simulate a region of 512 from the black hole, each in r and z direction, therefore the dimension of each cell is equivalent to 1. The gravity of the black hole is described by Paczyński and Wiita potential (Paczyński & Wiita 1980). In order to avoid the coordinate singularity on the horizon, the black hole is covered by a sink region of radius 3 around the origin, which do not affect the physics since the inner edge of KD is 3.
Citation: Roe 1981
Citation function: Uses
The CMEs are solar eruptions caused by magnetic reconnections in the star corona (Low 2001; Howard 2006). They expel a large amount of fast charged particles and a magnetic cloud that evolves into an interplanetary coronal mass ejection (ICME; Sheeley Jr. et al. 1985; Neugebauer & Goldstein 1997; Cane & Richardson 2003; Gosling 1990). If the ICME impacts the Earth, the measured SW dynamic pressure increases to 10−100 nPa and the IMF intensity to 100–300 nT (Gosling et al. 1991; Huttunen et al. 2002; Manchester IV et al. 2004; Schwenn et al. 2005; Riley 2012; Howard 2014; Mays et al. 2015; Kay et al. 2017; Savani et al. 2017; Salman et al. 2018; Kilpua et al. 2019; Hapgood 2019). The Disturbance Storm Time Index (Dst) indicates the magnetic activity derived from a network of near-equatorial geomagnetic observatories that measures the intensity of the globally symmetrical equatorial electrojet (the ring current), which is widely used to identify extreme SW and IMF space weather conditions (Sugiura & Chapman 1960; Loewe & Prölss 1997; Siscoe et al. 2006; Borovsky & Shprits, Yuri 2017). A negative Dst value means that Earth’s magnetic field is weakened due to the IMF erosion, particularly during solar storms. The strongest event observed so far is the Carrington event that occurred in 1859 (Carrington 1859). An unusual large number of sunspots on the solar disk and a wide active region was registered, and an extremely fast ICME was launched from it toward the Earth. Several authors studied the Carrington event and suggested that it was a shock that traveled at about 2000 km s−1 (Cliver et al. 1990) that generated the strongest geomagnetic storm with Dst ≈−1700 nT (Tsurutani et al. 2003). This was later revised to Dst ≈−850 nT by Siscoe et al. (2006). The most recent strongest event, called the Bastille Day event (14–16 July 2000), reached a Dst ≈−300 nT for an SW velocity of 1000 km s−1 and an IMF intensity of ≈45 nT (Rastatter et al. 2002). On the other hand, typical ICMEs impacting the Earth show an averaged plasma velocity of 350–500 km s−1 and IMF intensities between 9–13 nT, leading to geomagnetic storms with Dst  −50 nT (Cane & Richardson 2003).
Citation: Kay et al. 2017
Citation function: Background
Ring systems are a ubiquitous feature in planetary systems – all the gas giants in the Solar System have ring systems around them of varying optical depths (see, e.g., Tiscareno 2013; Charnoz et al. 2018), and ring systems have been detected around minor planets (e.g., Chariklo; Braga-Ribas et al. 2014), so it is reasonable that exoplanets and substellar objects host ring systems as well. Long-period eclipsing binary star systems, where one star is surrounded by an extended dark disk-like structure that periodically eclipses the other component, have already been observed, such as EE Cep (Mikolajewski & Graczyk 1999), ϵ Aurigae (Guinan & Dewarf 2002), and TYC 2505-672-1, with a companion period of 69 yr (Lipunov et al. 2016; Rodriguez et al. 2016). A large ring-like structure around a substellar companion was proposed to explain observations from 2007 from the J1407 system (Mamajek et al. 2012). 1SWASP J140747.93-394 542.6 (V1400 Cen; hereafter called “J1407”) is a young, pre-main-sequence star in the Sco-Cen OB association (Mamajek et al. 2012) with spectral type K5 IV(e) Li and is similar in size and mass to the Sun. In 2007, it displayed a complex symmetric dimming pattern of up to ~ 3 magnitudes during a 56 day eclipse. This has been attributed to the transit of a substellar companion (called “J1407 b”) with a mass of 60–100 MJup (Rieder & Kenworthy 2016) surrounded by an exoring system consisting of at least 37 rings and extending out to 0.6 au in radius (Kenworthy & Mamajek 2015). For these rings to show detectable transit signatures, they must be significantly misaligned with respect to the orbital plane of J1407 b (Zanazzi & Lai 2017). This potential ring system would be considerably larger than the ring system of Saturn, which is located within the planet’s tidal disruption radius. The proposed rings around J1407 b would even cover a significant fraction of the companion’s Hill sphere and would not be expected to be stable over gigayear timescales. If the candidate ringed companion is in a bound orbit around the star, this orbit must be moderately eccentric in order for no othereclipses to have been detected to date (Kenworthy et al. 2015), raising the possibility that there might be a second as yet undetected companion in the system that causes the implied orbital eccentricity for J1407 b. Radial velocity measurements are overwhelmed by the chromospheric noise of the star and do not place strong constraints on other substellar companions (Kenworthy et al. 2015). The transit of J1407 suggests that its orbital plane has a high inclination to our line of sight – if there are other planets inside the orbit of J1407 b, their orbits may well be coplanar with J1407 b and there is a high chance that these companions may transit J1407.
Citation: Rodriguez et al. 2016
Citation function: Background; Background
The X-ray luminosity of CG X-1 is variable by a factor of ≈10 (Bianchi et al. 2002; Weisskopf et al. 2004). The highest flux reported in the literature (5.2 × 10−12 erg cm−2 s−1 for a power-law fit or 5 × 10−12 erg cm−2 s−1 for an MCD fit, in the 0.5–8 keV band; Weisskopf et al. 2004) would imply, for a distance of 4.2 Mpc, a 0.5–10 keV luminosity of LX = (1.5–2) × 1040 erg s−1. If the system is Eddington-limited, the lower limit on the mass of the accreting BH is MBH ≳ 75 M⊙ for an He or C/O donor. For the system to shine in X-rays, the velocity of the WR star wind has to be slow enough to allow the formation of an accretion disc. This condition corresponds for CG X-1 to the requirement MBH ≳ 1.5 vw, 10004δ2 M⊙, where MBH is the BH mass, vw, 1000 is the wind velocity in units of 1000 km s−1, and δ ≈ 1 is a dimensionless parameter (adapted from Carpano et al. 2007b, see also Illarionov & Sunyaev 1975). In the simplest wind-accretion case (e.g. Edgar 2004), the luminosity can be estimated as

(1)

\begin{equation}
L_{\rm{X}}\approx \eta \frac{\dot{M}_{\rm{w}} c^2G^2M_{\rm{BH}}^2}{a^2 (v_{\rm{orb}}^2+v_{\rm{w}}^2)^2},
\end{equation}

where η is the efficiency, $\dot{M}_{\rm{w}}$ is the wind mass-loss rate, a is the orbital separation, vorb is the orbital velocity, and vw is the wind velocity at the BH orbit. Assuming $\dot{M}_{\rm{w}} = 10^{-5}$ M⊙ yr−1 and vw = 1000 km s−1 for the WR star (e.g. Crowther 2007), a = 5.8 × 1011 cm (for a 10 M⊙ companion), MBH = 75 M⊙, and the formation of a disc with η = 0.1, the corresponding luminosity is LX ≃ 2 × 1040 erg s−1. More in general, for MBH > 10 M⊙ and all the other things being equal, one finds LX ≳ 3 × 1039 erg s−1. In case of Roche lobe overflow, even higher X-ray luminosity could be achieved. However, we note that if CG X-1 is indeed a WR–BH binary, the WR star is probably not filling its Roche lobe [unless it is very massive; see for example the discussion of the case of Cyg X-3, where the orbital period is much shorter, in Szostek & Zdziarski (2008)]. An X-ray luminosity of ∼2 × 1040 erg s−1 can be therefore accounted for. We finally notice that, although we do not regard the question as crucial, the problem of the lifetime of the system discussed by Weisskopf et al. (2004) would be significantly attenuated, since the WR phase of a massive O-type star is thought to last a few ×105 yr (Meynet & Maeder 2005).
Citation: Carpano et al. 2007b
Citation function: Uses