adsabs/FOCAL huggingface dataset:
The wealth of precise all-sky data from the Gaia Data Release 2 (DR2) revealed a new feature in the Herzsprung–Russell diagram (HRD), namely a gap in the mid-M dwarf main sequence (Jao et al. 2018). The gap appears at a magnitude MG ∼ 10 and colour GBP – GRP ∼ 2.3−2.5 in the Gaia filter system. It is observed in optical and near-infrared colour-magnitude diagrams (CMDs), indicating that it is not specific to the Gaia photometry and not due to an atmospheric feature that would depend on the wavelength. Jao et al. (2018) suggest the feature is linked to the onset of full convection in M dwarfs. Interestingly, Mayne (2010) was the first to suggest the existence of an observational signature for the transition between fully and partly convective structures for pre-main sequence stars and predicted that it would result in a HRD gap. This author explored signatures of this transition in young clusters and linked the growth of a radiative core to rapid change in effective temperature caused by changes in the dominant energy transport mechanism and ignition of hydrogen burning. Following the announcement of the HRD gap discovery by Jao et al. (2018), MacDonald & Gizis (2018) proposed an explanation based on standard stellar evolution models. They suggest that the observed feature is due to the complex interplay between production of 3He and its transport by convection. More specifically, they predict a fast change in the luminosity over a narrow mass range of ∼0.31 M⊙−0.34 M⊙ characterised by the presence of a convective core and a convective envelope that ultimately merge. During this merging process, the central 3He abundance increases, causing an increase in luminosity and thus an observable feature in the luminosity function. In this paper, we re-examine the explanation suggested by MacDonald & Gizis (2018), since at first sight it is not clear why a sudden increase of the luminosity due to the merging of the convective zones of the core and the envelope would create a dip in the luminosity function, and thus a gap in the HRD. In this analysis, we confirm that the best explanation for the observed feature is linked to the property of 3He nuclear production and destruction, and to its mixing. We also find that a change in the energy transport from convection to radiation does not induce structural changes that could be visible. Regarding the very details of the process, however, we disagree with MacDonald & Gizis (2018) and propose an alternative explanation.
Citation: Mayne (2010)
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We used the combination of models from STARBURST99 and our model for the contribution from stripped stars to represent the radiation of a full stellar population in which stripped stars are formed. We made our models publicly available on the STARBURST99 online interface2, providing the addition from stripped stars to the spectral energy distribution, the emission rates of H I-, He I-, and He II-ionizing photons, and the high-resolution UV and optical spectra. In this study, we compare the contribution from stripped stars to the version of STARBURST99 that uses the initial mass function from Kroupa (2001) with mass limits of 0.1 M⊙ and 100 M⊙, that is, the same as what we employed for the stripped stars in the population. We chose to compare to the STARBURST99 models that use non-rotating stellar evolutionary models from the Geneva grids (Levesque et al. 2012) and spectral models from WM-BASIC for OB-stars (Pauldrach et al. 2001; Leitherer et al. 2010), CMFGEN for WR stars (Hillier & Miller 1998; Smith et al. 2002), and BASEL v3.1 for later type main sequence stars, cooler stars, and red supergiants (Lejeune et al. 1997). When comparing our models of various metallicity with STARBURST99, we used the STARBURST99 models of Z = 0.014, 0.008, 0.002, and 0.001 together with our models with metallicity of Z = 0.014, 0.006, 0.002, and 0.0002, respectively (see also Table A.1). We consider that the difference in metallicity between the two model sets is small and expect that the difference when using models with exactly the same metallicity is also small. We consider that the combination of the stripped star models and STARBURST99 is a good assumption for radiation with wavelengths shorter than ∼5000 Å. For the model to be accurate at longer wavelengths, we would need to decrease the radiation from giant stars to compensate for stars that we assume have become stripped. Giant stars emit their radiation primarily at wavelengths longer than ∼5000 Å. We expect the decrement of radiation at these long wavelengths to be at maximum about 30% (as also suggested by Eldridge et al. 2017, see e.g., their Figs. 5 and 15). This is approximately the fraction of massive stars that get stripped (Sana et al. 2012) and thus the fraction of giant stars that should be missing. This topic is beyond the scope of this paper, but we hope to address it more in detail at a later stage.
Citation: Leitherer et al. 2010
Citation function: Uses; Uses
Coming back to the flux values given in Table 1, it is tempting to explain the deficiency of the measured fluxes at energies ∼68 and ∼78 keV compared to the fluxes measured at 1.157 MeV as a consequence of an additional component of the flux at 1.157 MeV, which is generated by the interaction of low-energy cosmic rays (LECR) with abundant circumstellar species like Fe, Mn, and Cr in the environment of Cas A. Siegert et al. (2015) attempted to interpret these discrepant line fluxes of ${}^{44}_{}\mathrm{Ti}$  at different energies exactly in this way. However, considering that little or nothing is known about the LECR fluxes in the Cas A SNR and uncertainties in the generally small cross-sections for the relevant reactions to produce either ${}^{44}_{}\mathrm{Ti}$  or excited ${}^{44}_{}\mathrm{Ca}^{*}$ (Silberberg, Tsao & Barghouty 1998; Tsao, Silberberg & Barghouty 1998), the validity of such an interpretation of the discrepant measurements of ${}^{44}_{}\mathrm{Ti}$  fluxes at 68 and 78 keV and at 1.157 MeV by different instruments and by the same instrument (SPI) may be doubted. Additionally, if one assumes that the production of the 1.157 MeV line emission is supported in Cas A by the enhanced flux of LECRs, then immediately a problem arises related to the absence of other lines emission from much more abundant elements like oxygen, carbon, and nitrogen with very large, well-measured cross-sections for excitation. Indeed, simulations of the excitation emission lines produced by LECRs in Cas A have shown that lines at 4.43, 6.13, 6.9, and 7.1 MeV are expected to be strongest (Summa et al. 2011) for the expected spectrum of LECR (Berezhko et al. 2003) that are interacting with the CSM near the Cas A SNR. However, because no gamma-ray line emission at 4.43, 6.13, 6.9, and 7.1 MeV is observed from the Cas A SNR, we consider this explanation of the excitation origin of the high 1.157 MeV ${}^{44}_{}\mathrm{Ti}$  line flux compared to that at 68 and 78 keV as very problematic.
Citation: Tsao, Silberberg & Barghouty 1998
Citation function: Compare/Contrast; Compare/Contrast
Several evolutionary channels have been proposed that lead to a SNIa explosion. For a comprehensive review, see Livio & Mazzali (2018) and Wang (2018). Among these, the two classical scenarios are the single- and the double-degenerate channels. In the single-degenerate channel a white dwarf (WD) in a binary system accretes mass from a non-degenerate donor until it grows near the Chandrasekhar limit (Whelan & Iben 1973; Han & Podsiadlowski 2004; Nomoto & Leung 2018). In the double-degenerate channel two WDs in a close binary system merge due to angular momentum loss caused by the emission of gravitational waves (GWs) and the resulting merger has a mass near the Chandrasekhar limit (Whelan & Iben 1973; Iben & Tutukov 1984; Liu, Wang & Han 2018). Additional evolutionary channels for SNIa include the double-detonation mechanism (Woosley & Weaver 1986; Livne & Arnett 1995; Shen et al. 2012), the violent merger model (Pakmor et al. 2010; Sato et al. 2016), the core-degenerate channel (Sparks & Stecher 1974; Livio & Riess 2003; Kashi & Soker 2011; Wang et al. 2017) and a mechanism which involves the collision of two WDs (Benz, Thielemann & Hills 1989; Aznar-Siguán et al. 2013; Kushnir et al. 2013). In the double-detonation scenario a WD accumulates helium-rich material on its surface, which is compressed and ultimately detonates. The compression wave propagates towards the centre of the WD and a second detonation occurs near the centre of its carbon--oxygen core. In the violent merger model, the detonation of the WD core is initiated during the early stages of the merger. This can happen, for example, due to compressional heating by accretion from the disrupted secondary or due to a preceeding detonation of accreted helium (alike the double-detonation scenario) that is ignited dynamically (Guillochon et al. 2010; Pakmor et al. 2010, 2011, 2012, 2013; Kashyap et al. 2015; Sato et al. 2015, 2016). In the core-degenerate scenario a WD merges with the hot core of an asymptotic giant branch star during (or after) a common envelope (CE) phase. Finally, the evolutionary phase involving the collision of two WDs requires a tertiary star which brings the two WDs to collide due to the Kozai--Lidov mechanism, or dynamical interactions in a dense stellar system, where this kind of interaction is more likely to happen.
Citation: Kashi & Soker 2011
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Application of linear theory to investigate IGWs has been done extensively in the field of oceanography (Sutherland 2010). The general idea is to linearize hydrodynamical equations and introduce a wave-like ansatz to solve the linearized equations. In the context of stellar parameters, one of the earliest works was done by Press (1981), which focused on solar-like stars. In our work, we follow similar steps, starting with the linearized hydrodynamical equations in the anelastic approximation:
(8)$$\begin{eqnarray*}
\nabla \cdot \overline{\rho } \boldsymbol {v} = 0,
\end{eqnarray*}$$(9)$$\begin{eqnarray*}
\frac{\partial \boldsymbol {v}}{\partial t} = - \nabla \left(\frac{P}{\overline{\rho }}\right) - C \overline{g} \boldsymbol {\hat{r}}
\end{eqnarray*}$$(10)$$\begin{eqnarray*}
\frac{\partial T}{\partial t} = - v_\mathrm{ r} \left(\frac{\partial \overline{T}}{\partial r} - (\gamma - 1) \overline{T} h_\rho \right)
\end{eqnarray*}$$We ignore the pressure term in equation (4), which leads to a set of equations that conserve energy (Brown, Vasil & Zweibel 2012). We also do not consider rotational, thermal diffusion, or viscous effects. In cylindrical coordinates, this 2D analysis allows us to set the z-derivatives to zero. The three equations shown above can then be reduced to one second-order differential equation with vr(r) as the evolving term:
(11)$$\begin{eqnarray*}
0 &amp;=&amp; \frac{\partial ^2 \alpha }{\partial r^2} + \left(\frac{N^2}{\omega ^2} - 1 \right) \frac{m^2}{r^2} \alpha + \left[ -\overline{\rho }^{-1/2}\frac{\partial ^2 \left(\overline{\rho }^{1/2}\right)}{\partial r^2} + \frac{\partial h_{\rho }}{\partial r} \right]\alpha \nonumber \\
&amp;&amp;+\, \frac{1}{4r^2} \alpha,
\end{eqnarray*}$$where $\alpha = v_\mathrm{ r} \overline{\rho }^{1/2} r^{3/2}$. We have used a wave ansatz of the form vr(r, θ, z) ∝ vr(r)eimθe−iωt, where m is the horizontal wavenumber,1 ω is the angular frequency, and θ is the angular coordinate in the equatorial plane. Generally, in the radiation zone, there will be regions where the oscillatory term (OT), (N2/ω2 − 1)m2/r2, dominates and regions where the density term (DT), $[ -\overline{\rho }^{-1/2}\partial ^2 (\overline{\rho }^{1/2})/\partial r^2 + \partial h_{\rho }/\partial r]$, dominates. When the ratio of OT to DT is less than 1, an IGW loses its wave-like behaviour and the approximate radius where this ratio is exactly equal to 1 is called the turning point. The importance of this will be discussed in Section 4.
Citation: Brown, Vasil & Zweibel 2012
Citation function: Uses
There are several different effects that have been considered to explain the observed offset in the [O iii]λ5007/Hβ vs. [N ii]λ6584/Hα BPT diagram among high-redshift galaxies. First, there are the physical parameters describing the H ii regions contributing to the integrated line ratios from galaxies. These include the ionization parameter, ionizing spectrum of the stars illuminating the H ii-region gas, and the electron density. Early work highlighting the issue of the high-redshift offset in the [O iii]λ5007/Hβ vs. [N ii]λ6584/Hα BPT diagram focused on these parameters (e.g., Shapley et al. 2005; Brinchmann et al. 2008; Liu et al. 2008), and they have been revisited more recently by Kewley et al. (2013), Steidel et al. (2014), and Masters et al. (2014). Systematically higher ionization parameters (which appear to apply in high-redshift galaxies; Nakajima et al. 2013), harder ionizing spectra, and higher electron densities (Shirazi et al. 2014), all tend to shift the locus of galaxies in the [O iii]λ5007/Hβ vs. [N ii]λ6584/Hα diagram toward higher [O iii]/Hβ and [N ii]/Hα values. Next, there is the role of different types of pressure in determining the internal structure and dynamics of H ii regions. Yeh et al. (2013) and Verdolini et al. (2013) suggest that radiation pressure is significant in high-redshift H ii regions, as compared with gas pressure associated with stellar winds, and that the effects of radiation pressure can lead to [O iii]/Hβ line ratios in excess of the “maximum starburst” limit of Kewley et al. (2001). Possible contamination by weak AGNs has also been suggested as a way to shift galaxy emission-line ratios into the “composite” region of the [O iii]λ5007/Hβ vs. [N ii]λ6584/Hα BPT diagram (e.g., Wright et al. 2010), in between the curves of Kauffmann et al. (2003) and Kewley et al. (2001). Finally, both Masters et al. (2014) and Steidel et al. (2014) consider gas-phase abundance ratios—specifically, the N/O ratio—which can affect where galaxies fall in the [O iii]λ5007/Hβ vs. [N ii]λ6584/Hα BPT diagram. If the relationship between N/O and 





 evolves out to high redshift, then distant galaxies will shift relative to local ones in the [O iii]λ5007/Hβ vs. [N ii]λ6584/Hα BPT diagram.
Citation: Liu et al. 2008
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Extensive modelling efforts have been undertaken to explain observations and the formation of dust in local and high redshift SNe. While some studies use the classical nucleation theory (CNT) to describe dust formation under equilibrium conditions in shocked environments (Kozasa et al. 1989; Todini & Ferrara 2001; Nozawa et al. 2003; Schneider et al. 2004), other studies describe the formation of both molecules and dust clusters from the shocked gas by assuming a chemical kinetic approach under non-equilibrium conditions and the subsequent coalescence and coagulation of these clusters to form dust grains (Cherchneff & Lilly 2008; Cherchneff & Dwek 2009; 2010; Sarangi & Cherchneff 2013; 2015, hereafter SC15). On the other hand, the processing of dust in SNRs has received less attention. The thermal sputtering of dust by the forward shock has been studied by Nozawa et al. (2006), and sputtering by the RS in local and primitive SNe was modelled by Nozawa et al. (2007), Bianchi & Schneider (2007), and Nath et al. (2008). Sputtering by the RS in Type II-b SNe with application to the Cas A SNR was studied by Nozawa et al. (2010). Finally, Silvia et al. (2010; 2012) proposed a hydrodynamic study of an ejecta clump, which was processed by the RS, and followed thermal sputtering as the clump was crossed and gradually disrupted by the shock. All studies show that dust is strongly reprocessed in SNRs. However, the final mass of surviving dust is not well constrained because of several factors. Firstly, all studies use as initial conditions for sputtering in the SNR the dust masses and size distributions that result from applying CNT. The chemical types and masses of the dust that are used as initial conditions may consequently be in error using this formalism, as discussed by Cherchneff (2014). Furthermore, some of these studies assume a homogeneous ejecta in the SNR phase (Nozawa et al. 2007; 2010; Bianchi & Schneider 2007; Nath 2008), which implies that very harsh conditions hold in the post-shock gas for dust survival because the RS crosses and reprocesses the ejecta at very high velocities. By considering clumpy ejecta, the RS velocity will be greatly reduced in the dense clumps (Silvia et al. 2010; Biscaro & Cherchneff 2014).
Citation: Sarangi & Cherchneff 2013
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The planetary physical properties that can be derived from measurements with the two most successful exoplanet detection techniques, that is, with the radial velocity method and the transit method, are the planet radius (if the planet transits its star), mass (a lower limit if the planet does not transit its star), and the orbital period and distance, eccentricity, and inclination (if the planet transits its star) (see Perryman 2018, and references therein). The upper atmospheres (above optically thick clouds) of giant exoplanets and those of smaller planets around M stars can be probed using spectroscopy during primary and/or secondary transits (Bean et al. 2010; Swain et al. 2009, 2008; Tinetti et al. 2007; Nakajima 1983). It appears to be virtually impossible to characterize the (lower) atmospheresand surfaces of small, Earth-like planets in the habitable zones of solar-type stars (Bétrémieux & Kaltenegger 2014; Misra et al. 2014; Kaltenegger & Traub 2009), in particular because the light of the parent star is refracted while traveling through the lower atmosphere of its planet and emerges forever out of reach of terrestrial telescopes (García Muñoz et al. 2012). The (lower) atmosphere and surface of a planet are crucial for determining the habitability of a planet, as they hold information about cloud composition, trace gases in disequilibrium and probably most importantly, liquid surface water (see, e.g., Schwieterman et al. 2018; Kiang et al. 2007a,b, and references therein). For such a characterization of terrestrial-type planets, direct observations of the thermal radiation that they emit or of the light of their parent star that they reflect are required. The numerical results that we present in this paper concern the reflected starlight. Because of the huge distances involved, any measured reflected starlight pertains to the (illuminated and visible part of the) planetary disk. It therefore is a disk-integrated signal.
Citation: Swain et al. 2009
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Citation: Nakajima 1983
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Information about the composition and structure of astrophysical ices, and in particular the presence of solid methanol, is obtained by comparing astrophysical observations in the infrared to laboratory data (Dartois et al. 1999; Pontoppidan et al. 2003; Boogert et al. 2008; Bottinelli et al. 2010). The main infrared absorption features used for solid CH3OH identification in interstellar ices are the 3.53 μm (2833 cm−1) and 9.74 μm (1027 cm−1) bands. They correspond to the CH3 symmetric stretching mode and the CO
stretching mode, respectively. The laboratory data most widely used for comparison to observations are those provided by d’Hendecourt & Allamandola (1986), Hudgins et al. (1993), and Kerkhof et al. (1999). The first reference gives IR spectra and band strengths of methanol ice at 10 K in the 2.5–20 μm range. The second reference presents a thorough infrared study providing a compendium of band strengths and optical constants in the 2.5–200 μm interval for ices of astrophysical relevance. For pure methanol, optical constants are given for an ice grown at 10 K and subsequently warmed to 50 K, 75 K, 100 K, and 120 K. Band strengths at 10 K are also given. As far as we know, these are the only optical constants available in the literature for solid methanol. The third reference reports band strength variations of methanol absorptions in the mid-IR region with dilution in H2O ice, or in H2O and CO2 ice mixtures. Additional laboratory works provide infrared band strengths of this species at 10 K. Specifically, the data in Palumbo et al. (1999) correspond to the mid-IR (MIR), and the data in Sandford & Allamandola (1993) and Gerakines et al. (2005) to the near-IR (NIR). Finally, a recent paper by Bouilloud et al. (2015) presents a compilation and new measurements of MIR band strengths of methanol ice at low temperatures. Since methanol ice densities were unknown, the uncertainty on band-strength determination is very large. In particular, a discrepancy among the literature results up to 40% for the strength of the 9.74 μm band has been pointed out.
Citation: d’Hendecourt & Allamandola (1986)
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In general, the most important effect of the ICME passage in generating major FDs is ascribable to the energetic interplanetary shock/sheath region, although it has been proposed that the MC effect could even be dominant (e.g., Sanderson et al. 1990). Intense shocks are associated with the fast ICME propagation and may overcome the MC effect in modulating the GCR intensity, as turbulent magnetic fluctuations within the sheath can influence the GCR propagation significantly (e.g., Wibberenz et al. 1998). However, the minimum intensity occurs after the arrival of the MC (Badruddin 1986; Zhang & Burlaga 1988). For slower ICMEs, the effect of shock and turbulent sheath on the GCR intensity could be similar to that associated with the MC and sometimes it could be negligible. Therefore, if no interplanetary shock is found leading an ICME and a weak magnetic field turbulence is observed in the sheath region, the FD evolution mostly depends on the magnetic configuration of the ICME. A statistical study was performed by Richardson & Cane (2011) on more than 300 ICMEs, showing that 80% of them is associated with an FD, detected by the anticoincidence guard data on the International Monitoring Platform (IMP-8; Cane 2000), and that the minimum GCR intensities occur within the MC. They also concluded that ICMEs containing MCs cause deeper FDs on average with respect to ICMEs that do not have any MC structure. Despite numerous attempts to relate the properties of FDs with those of ICMEs at 1 au (Richardson et al. 1996; Belov et al. 2001, 2014; Belov 2008; Dumbović et al. 2012) there are significant gaps in our understanding of their underlying physical mechanisms. A renewed interest in studying GCR FDs has been fostered by the most recent S/C observations in the near-Earth space at the lower rigidities with respect to NMs, for which FDs appear larger allowing for the study of the fine structure of the decrease formation. For instance, different data sets have been used to study GCRs and the associated FDs recorded by a radiation monitor on board LISA Pathfinder (LPF; Armano et al. 2018), the Electron Proton Helium INstrument detector on board the Solar and Heliospheric Observatory (SoHO) and the Chandra X-ray observatory (Heber et al. 2015; Dumbović et al. 2018), and the anticoincidence shield of the International Gamma-Ray Astrophysics Laboratory’s spectrometer (Jordan et al. 2011). FDs have also been observed at several solar distances by the Radiation Assessment Dosimetry dose rates on the Mars Science Laboratory (Guo et al. 2018; von Forstner et al. 2018, 2020; Papaioannou et al. 2019), the Cassini’s Magnetosphere Imaging Instrument and Low Energy Magnetospheric Measurement System measurements (Roussos et al. 2018) and beyond (Witasse et al. 2017; Winslow et al. 2018).
Citation: Cane 2000
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