adsabs/FOCAL huggingface dataset:
Major differences between our approach and a full re-calculation of the hydrodynamics were not expected since the changes in the 14N and 22Ne abundances at the different metallicities do not alter the energy release in the hydrodynamic simulations significantly. The situation is different for deflagrations where the buoyancy, and therefore the Rayleigh-Taylor instabilities, depend on Ye. In contrast to detonations, the propagation of a deflagration front is thus affected by the metallicity (e.g., Meakin et al. 2009). The assumption we made here is confirmed by the comparison of the models presented in Table 1. Model M2a is taken from Gronow et al. (2020). The model was calculated at zero metallicity and has a total mass of 1.05 M⊙ with a He shell of 0.07 M⊙ at He ignition. Model M10_05_1, on the other hand, has a similar mass configuration, though it was calculated at solar metallicity (Model M10_05 in Gronow et al. 2021). Model M2a_pp is the same model as Model M2a, but the postprocessing step was calculated with solar metallicity instead of zero metallicity. An inspection of the abundances of Models M2a_pp and M10_05_1 at t = 100 s after He detonation ignition shows that the results of the postprocessing step with varying metallicities are in reasonably good agreement with a full re-calculation of the hydrodynamic model. The maximum difference in the yields produced in the core detonation is only 10%, while the maximum difference is 50% in the He detonation (excluding 12C in both). However, differences in the yields produced in the He detonation can in part be attributed to the slightly different setups of Model M2a (and therefore Model M2a_pp) and Model M10_05_1 at the beginning of the relaxation simulation, with the differences in the total and shell masses being less than 1% (see Gronow et al. 2021 for an explanation of the difference). Generally, the contribution of the yields from the He detonation to the total nucleosynthetic yields are small compared to those of the core detonation. Our approach is thus sufficient to derive nucleosynthetic yields for GCE studies. It saves significant computational costs as additional 3D hydrodynamical simulations of the explosion do not need to be carried out. Nevertheless, there might be slight differences visible in the observables because they are sensitive to the products of the He shell detonation (Höflich et al. 1996; Nugent et al. 1997; Kromer et al. 2010).
Citation: Gronow et al. 2021
Citation function: Differences; Differences
The objective of this article is to study the spatial distribution and possible changes in the properties of carbon nano-dust in protoplanetary disks (PPDs). Carbon nanodust, detected under more or less organised structures and different ionisation states, constitutes a major component of dust in the interstellar and circumstellar environments. Vibrational emission bands in the near- to mid-IR from nanocarbon dust have been observed towards PPDs around most of the Herbig Ae stars, about half of the Herbig Be stars, and a few T-Tauri stars (e.g. Brooke et al. 1993; Acke & van den Ancker 2004; Acke et al. 2010; Seok & Li 2017). In contrast to large grains, these tiny and numerous carbon grains are well coupled to the gas and do not settle towards disk midplanes. This results in different spatial distributions in which tiny grains are present at the disk surfaces (e.g. Meeus et al. 2001; Habart et al. 2004; Lagage et al. 2006) and in the cavity or gaps from which the pebbles are missing (e.g. Geers et al. 2007; Kraus et al. 2013; Klarmann et al. 2017; Kluska et al. 2018; Maaskant et al. 2013). The very small carbon grains in the irradiated disk layers may have strong consequences (e.g. Gorti & Hollenbach 2008). As in the irradiated regions of the interstellar medium, they are the prevalent contributors to the energetic balance because they are very efficient at absorbing UV photons and heating the gas through the photoelectric effect. The highest fluxes of lines tracing the warm gas (e.g. [OI] 63 and [OI] 145 μm, H2 0-0 S(1), and high-J CO) are found in PPDs that show a large amount of flaring and high aromatic band strength (e.g. Meeus et al. 2013). Moreover, due to their large effective surface area, they may dominate the catalytic formation of key molecules as H2 and the charge balance. The disk structure may further depend on the level of nanograins that are coupled with the gas. Characterising the size and properties of these tiny grains through the disks, from internal to external regions, is thus of prime importance to understand the structure and evolution of PPDs.
Citation: Geers et al. 2007
Citation function: Compare/Contrast; Compare/Contrast
Our first attempt to derive the $PL_{K_\mathrm{s}}$ relation has revealed a number of RRLs brighter than the main relation. We investigated their spatial distribution and found that they are mainly located in the central part of the LMC. This is shown in Fig. 5 which presents in the bottom panel the spatial distribution of the whole sample of RRLs considered in this paper (${\sim}29\, 000$), in the top panel the distribution of the RRLs which appear to be overluminous in the $PL_{K_\mathrm{s}}$ relation, and, in the middle panel, the distribution of the RRLs which were actually used to fit our final $PL_{K_\mathrm{s}}$ relation. Further investigations, performed using the Fourier parameters (ϕ31, ϕ21, R21, and R31) of the light curves available in the OGLE IV catalogue did not show any particular properties of the overluminous RRLs. On the other hand, in the period–amplitude diagram based on the I amplitudes available in the OGLE IV catalogue, the bright RRLs all show small, in some cases near-zero amplitudes, compared with regular RRLs of the same period. The decrease in amplitude at a given period can be owing to these RRLs being blended with non-variable stars. We expect the centroid of a blended source to be determined with poor accuracy (see e.g. Ripepi et al. 2014, 2015). For this reason as a further test we plotted the distribution of distances in arcsec of the VMC sources cross-matched with the OGLE IV RRLs. A clear separation is now seen in the two samples, for 94 per cent of the RRLs lying on the PL/PW relations the cross-match radius is less than 0.2 arcsec, whereas for 68 per cent of the overluminous RRLs the cross-match radius is larger than 0.2 arcsec. The average accuracy of the VMC astrometry is of the order of 0.080 arcsec both in RA and in Dec. (Cioni et al. 2011). We therefore discarded the RRLs with a cross-match radius larger than 0.2 arcsec. A total of 3252 objects were discarded. This procedure allowed us to significantly reduce the scatter on the PL/PW relations. We visually inspected the VMC images of some of the discarded RRLs, confirming that they all are clearly blended with stars and/or background galaxies. Similar investigations were performed in the past and the same effects were noted, e.g. by Ripepi et al. (2015) for Type II Cepheids. The final sample of LMC RRLs after this cleaning procedure contains 25 795 objects. This is the sample that was used as a starting point to investigate the PL and PW relations presented in the following sections. Additional RRLs were later discarded from the final fit based on a 3σ clipping procedure leading to $PL_{K_\mathrm{s}}$ relations using a clean sample of ${\sim}22\, 000$ RRLs.
Citation: Ripepi et al. 2014
Citation function: Uses; Uses
In direct relation to the interplay between the structure of the dark matter distribution and the baryon physics, galaxies are found in a wide range of structural hierarchies, from low-density regions to groups and clusters (see e.g. Tully 1987; Berlind et al. 2006; Yang et al. 2007), and during their lifetime they experience merging events (e.g. Mamon 1988; Tempel et al. 2017). In some cases, the mergers eventually devoid their entire neighbourhood, leaving behind a single elliptical galaxy of group-scale mass, called a fossil group galaxy (Ponman et al. 1994; Jones et al. 2003). Numerical simulations by Barnes (1989) first motivated such a hierarchical merging scenario (see also Díaz-Giménez, Muriel & Mendes de Oliveira 2008). Since then, there have been several supporting reports of X-ray sources identified as fossil groups (Santos, Mendes de Oliveira & Sodré 2007; La Barbera et al. 2009). Because most fossil systems found to date lie within z  0.2, fossil galaxy groups most likely are old, undisturbed systems due to the lack of major mergers. While some luminous galaxies experience major merger events in their evolution, fossil group galaxies acquire their mass typically through minor merger events, where the mass ratio stays below 0.3 (DOnghia et al. 2005). Simulations show that a fossil system may assemble half of its mass in dark matter by redshift z > 1, and that the assembled mass at any redshift is generally higher in a fossil than in regular groups (Dariush et al. 2007). Since this merging process is relatively fast compared to the cooling time of the surrounding gas, comparable to one to several Hubble times, fossil groups are usually found embedded in giant, X-ray luminous gas haloes (Mulchaey 2000). If a fossil system has not yet fully merged, it can be identified by another criterion, a gap in brightness of at least 2.5 mag (usually defined in the r band) between the first and fourth brightest galaxies in the group (Dariush et al. 2010; Zarattini et al. 2014) within half its virial radius.
Citation: Jones et al. 2003
Citation function: Background
At present, however, kinematical studies of the nearby late-type stars have become somewhat foreshadowed with the advent of deep, wide-field SDSS (York et al. 2000), SDSS/SEGUE (Yanny et al. 2009) and RAVE (Steinmetz et al. 2006) surveys which have made it possible to study K–M dwarfs to much greater distances from the Sun. Using their spectroscopic catalog with radial velocities measured with an external accuracy of 7–10 km s-1 and applying photometric parallax relations, the SDSS teams have traced K–M dwarfs in the distance range up to ~2 kpc, thus providing valuable information on the spatial, velocity, and metallicity distributions with respect to vertical distance from the Galactic plane (Bochanski et al. 2007, 2011; Jurić et al. 2008; Fuchs et al. 2009; Bond et al. 2010; West et al. 2011; Schlesinger et al. 2012; Zhang et al. 2013), as well as on the luminosity and mass functions of low-mass dwarfs in the Galactic disk (Covey et al. 2008; Bochanski et al. 2010). Relatively local (| z |   500 pc) samples of late-type dwarfs from the RAVE survey, based on radial velocities accurate to ~2 km s-1 and photometrically determined distances, have been used to deduce the solar space velocity with respect to the Local Standard of Rest (Coşkunoǧlu et al. 2011; Pasetto et al. 2012a; Golubov et al. 2013). The RAVE survey has also resulted in a catalog of ~44 000 candidate active stars (Žerjal et al. 2013) which makes a major contribution to the data on chromospheric emission of cool dwarfs. In the context of these massive surveys, as well as of present-day models of the Galaxy, the samples of nearby late-type dwarfs with most accurate trigonometric parallax and radial-velocity measurements, as well as other high-quality observational data, still remain important, as they provide a fundamental framework for calibration and tests of relations between low-mass star parameters (such as, e.g., color-luminosity, mass-luminosity, activity-age, chemo-kinematic relations).
Citation: Bochanski et al. 2010
Citation function: Background; Background
There are several reasons to suspect that the (galaxy-sourced) ionizing background during reionization may have been somewhat harder than the estimates given here. First, our calculations neglect the filtering effects of optically thick H i in the ISM of the host galaxy, and within the cosmic web. Absorption by this gas would have hardened the spectrum of the ionizing radiation as it escaped the galaxy and traveled through the IGM (e.g., Madau 1995; Faucher-Giguère et al. 2009; Haardt & Madau 2012). These effects were likely strongest during the last stages of reionization, when the radiation typically had to travel large distances to reach the I-fronts. Second, our calculations neglect the effects of binary star systems. Mass transfers and mergers between binary companions can extend the period over which ionizing photons are produced by the stellar population, which would harden the time-integrated spectrum (Eldridge & Stanway 2009; Stanway et al. 2016). Lastly, recent studies have suggested that the IMF in starburst galaxies may be more top-heavy than the IMF assumed here (Baugh et al. 2005; Gunawardhana et al. 2011; Marks et al. 2012; Zhang et al. 2018). Most recently, Schneider et al. (2018) measured a logarithmic slope of 





 in the mass range 15–200 M⊙, using spectroscopic measurements of the 30 Doradus star-forming region in the Large Magellanic Cloud. (The IMF adopted here has a slope of 2.3 for M > M⊙, and a cutoff of 120 M⊙.) Each of the provided effects would work in the direction of making 





 smaller. Based on these considerations, we argue that the lower half of Figure 2, with 





 ≲ 1.5, is likely the most relevant region of parameter space for 





. In what follows, we adopt 





 = 1.5 as our fiducial value, but we note that 





 is only mildly sensitive to 





 except at the fastest I-front speeds. In the next section we find that 





 = 104 km s−1 is close to the upper limit achieved by I-fronts in cosmological simulations, which yields 





 = 26,200 K, assuming 





 = 1.5 (see Figure 2). This result varies by Δ





 = −4200 (+1600) K if we instead assume 





 = 2.5(0.5).
Citation: Stanway et al. 2016
Citation function: Compare/Contrast
Recently, the EHT Collaboration completed an analysis of the linear polarization of M87*, providing novel insights into its magnetic field structure in particular (Event Horizon Telescope Collaboration 2021a,b; Goddi et al. 2021). Synchrotron emission, which dominates the millimetre image, is initially generated perpendicular to the local magnetic field, such that its polarization carries with it an imprint of the field geometry (Palumbo, Wong & Prather 2020). Then, as this polarization propagates, it is further modified by Faraday effects, important for depolarizing accretion flows down to observed levels and generating the observed rotation measure (Ballantyne, Özel & Psaltis 2007; Mościbrodzka et al. 2017; Jiménez-Rosales & Dexter 2018; Ricarte et al. 2020). Fully polarized radiative transport simulations on EHT scales have been developed in the past few decades, allowing us to link polarized images to the detailed physics of the underlying plasma and the space–time producing them (Bromley, Melia & Liu 2001; Broderick & Loeb 2006; Broderick & McKinney 2010; Porth et al. 2011; Shcherbakov, Penna & McKinney 2012; Dexter 2016; Mościbrodzka & Gammie 2018). This enabled (Event Horizon Telescope Collaboration 2021b) to discriminate between two major classes of accretion disc: a ‘Magnetically Arrested Disk’ (MAD) and ‘Standard and Normal Evolution’ (SANE). MAD accretion discs have magnetic fields strong enough to affect the disc dynamics and exhibit stronger poloidal (or non-toroidal) magnetic field components (Bisnovatyi-Kogan & Ruzmaikin 1974; Igumenshchev, Narayan & Abramowicz 2003; Narayan, Igumenshchev & Abramowicz 2003; Chael, Narayan & Johnson 2019). Meanwhile, the weaker magnetic fields of a SANE disc are sheared out by the motion of the plasma into a mostly toroidal configuration (Narayan et al. 2012; Sądowski et al. 2013; Ryan et al. 2018). The fractional linear polarization of M87*, the upper limit on its circular polarization, and most importantly the ‘twisty pattern’ of its spatially resolved linear polarization map favour ‘MAD’ models of M87* (Event Horizon Telescope Collaboration 2021b).
Citation: Mościbrodzka et al. 2017
Citation function: Background
CMEs are known for large-scale expulsion of magnetized plasma structures from closed magnetic field regions on the Sun. They were first detected in the coronagraphic images taken in 1971 by NASA’s OSO-7 spacecraft (Tousey 1973). However, some definite inferences for the solar wind (Eddington 1910; Birkeland 1916; Biermann 1951) as well as CMEs from the Sun (Chapman & Ferraro 1931; Eddy 1974) were made decades before their formal discovery. Following OSO-7, a series of spacecraft (Skylab, Helios, P78-1 Solwind, SOHO, Coriolis, and STEREO, etc.) have observed thousands of CMEs leading to a vast literature (Munro et al. 1979; Howard et al. 1985; Gosling 1993; Hundhausen 1999; Gopalswamy et al. 2000; Schwenn 2006; Vourlidas et al. 2010; Chen 2011; Wang et al. 2011; Webb & Howard 2012; Mishra & Srivastava 2013; Mishra et al. 2017; Harrison et al. 2018). CMEs have been observed to occur often having spatial and temporal relation with solar flares, eruptive prominences (Munro et al. 1979; Webb & Hundhausen 1987; Zhang et al. 2001; Gopalswamy et al. 2003) and with helmet streamer disruptions (Dryer 1996). Unlike CMEs from the Sun, to observe stellar CMEs are challenging because the close stellar environment cannot be spatially resolved. Although stellar CMEs have not yet been directly detected in Thomson-scattered optical light from other stars, it is believed that the extreme X-ray flares observed on stars may be in conjunction with extreme stellar CMEs (Houdebine, Foing & Rodono 1990; Wheatley 1998; Leitzinger et al. 2011; Aarnio, Matt & Stassun 2012; Osten & Wolk 2015; Vida et al. 2016). Indeed, the stellar X-ray flare, helmet streamers, and prominences observed on T Tauri Stars have shown similarities with those observed on the Sun (Haisch, Antunes & Schmitt 1995; Massi et al. 2008). The CMEs and flares themselves may not be causally related, they both seem to be involved with the reconfiguration of complex magnetic field lines within the corona caused by the same underlying physical processes, e.g. magnetic reconnection (Priest & Forbes 2002; Compagnino, Romano & Zuccarello 2017). But, even for the sun, it has been noted that not all flares are accompanied by CMEs and not all CMEs by flares (Munro et al. 1979; Harrison 1995; Yashiro et al. 2008b; Wang & Zhang 2008).
Citation: Webb & Hundhausen 1987
Citation function: Background
Citation: Schwenn 2006
Citation function: Background
The first INM of PG 1553+113 to our knowledge was reported by Stalin et al. (2005) – they applied the C-test to the R-band LCs (total duration of about 10 h) and found the source to be variable in one night out of two. Another two nights of INM about 13 h long were presented by Osterman et al. (2006). They found no significant INV, but they did not present statistical tests. To quantify the conclusion of Osterman et al. (2006), we used the C-test in the form C = σ/⟨e⟩, where σ is the standard deviation of the source LC and ⟨e⟩ the mean uncertainty of the source photometric measurements. We applied this test to the data listed in Table 4 of Osterman et al. (2006) and found C = 1.52 and C = 1.28 for the respective nights, which means that PG 1553+113 was non-variable. Andruchow et al. (2007, 2011) reported the results from the INM they carried out during four nights in the VR-bands (April 21 to 24, 2007) and five nights in the BR-bands (April 21 to 25, 2009), respectively. Applying the C-test, the authors found no significant variability. Gopal-Krishna et al. (2011) detected INV during three nights out of three using C- and F-tests (a total monitoring duration of 16 h in the R-band). Gaur et al. (2012) found no INV during six nights of monitoring in the BR-bands (C- and F-tests were applied). Gupta et al. (2016) presented the results from the INM during seven nights in the R-band (a total monitoring duration of about 26 h). The authors found the source to be variable in one night, non-variable in another, and probably variable in the remaining nights (F- and χ2 tests were applied). We, however, should point out that in the latter nights the variability amplitude seems to be quite close to the magnitude uncertainties. In such cases the usage of the C-test could be more appropriate (see Zibecchi et al. 2017, 2020). Pandey et al. (2019) monitored PG 1553+113 for eight nights in the VR-bands for 2–4 h each night. Employing enhanced F- and nested ANOVA tests, the authors found the source to vary on intra-night timescales for three nights. Finally, Pasierb et al. (2020) found no INV in the BVR-bands during a single night of monitoring (duration of about 3.7 h, F-test employed).
Citation: Gopal-Krishna et al. (2011)
Citation function: Compare/Contrast
In order to assess the impact of the brightness enhancement due to the phase function effects in the light scattering of the coma dust, we apply different phase function corrections to the measured coma brightness of 67P for the 10 000 km aperture radius. In general, the phase function reduces the reflected coma magnitude value. For small phase angles α (typically up to 30°), a linear reduction with α can be applied with typical parameters β for comets between 0.02 and 0.06 (Lamy et al. 2004). Alternatively, the two-parameter phase function of Schleicher, Millis & Birch (1998), based on 1P/Halley measurements, is available or the phase angle correction by Müller (1999) that also considers geometric projection effects by a non-spherical coma, symmetric to the radial direction with respect to the Sun. Fornasier et al. (2015) obtained results for the 67P nucleus surface reflectivity applying a HG-type phase function (Bowell et al. 1989) with a best-fitting G of −0.13. Fig. 5 shows the phase angle-corrected afρ values of 67P for the various phase function solutions. In Fig. 5, β = 0.04 is used as mean value for the linear phase function that is lower than the β value found for the nucleus of 67P (0.059–0.076; Tubiana et al. 2011; Lowry et al. 2012). A second HG-type phase function is applied as well with G = 0.15 that is found from many asteroid light curves (Tedesco 1989). Obviously, the afρ value for zero phase angle depends on the phase darkening and is highest for the Bowell-type phase function with G = −0.13 (as obtained for the 67P nucleus by Fornasier et al. 2015). Two phase functions provide a close to linear decrease of afρ with solar distance over the time interval from the end of 2015 September to the beginning of 2016 April, i.e. the linear phase function with β = 0.04 and the Bowell-HG-type phase function with G = 0.15. The phase functions of Schleicher et al. (1998) and of Müller (1999) can approximate the afρ data well between the end of September and end of December, but do not describe well the 67P dust activity measured in 2016.
Citation: Bowell et al. 1989
Citation function: Uses; Uses