adsabs/FOCAL huggingface dataset:
In the southern sky, the large background of atmospheric muons reduces the efficiency to select through-going tracks induced by neutrinos below the PeV regime. A very large fraction of the aforementioned background can be rejected by imposing an active veto at the detector boundary, as, for example, used in Aartsen et al. (2013a). This reduces the detector volume to a smaller fraction of the instrumented volume sacrificing statistics for signal purity. Furthermore, the more clearly an event is identified as a starting track, the more probable it is to be an astrophysical rather than atmospheric background. Down-going atmospheric neutrino events at high energy are likely to be accompanied by muons produced in the same cosmic-ray shower that triggers the veto and reduces the atmospheric neutrino background (Schönert et al. 2009; Gaisser et al. 2014). In analyses using veto techniques (Aartsen et al. 2013a, 2015b), the selection is usually more efficient for cascade-like events than tracks, and high astrophysical neutrino purity demands neglecting energies below 60 TeV, where backgrounds are more abundant. In searches for point-like sources of astrophysical neutrinos, track-like events are of great importance given their good angular resolution compared to cascade-like events. Furthermore, the purity demands are lower since the signal of a point-like source is restricted to a small portion of the sky, hence reducing the background significantly. Consequently, the minimum required total charge deposited in the PMTs of the IceCube detector by an event is lowered to 1500 p.e. compared to 6000 p.e. (Aartsen et al. 2013a), resulting in a higher signal efficiency at lower energies. In addition, only down-going tracks are used, and cuts are imposed that select well-reconstructed track-like events (Aartsen et al. 2016a). For 





 events at energies smaller than 200 TeV, the effective area of the analysis is bigger than for vertically through-going tracks (δ  −30°; Figure 1). For energies up to 1 PeV, the effective area is smaller, but a higher purity is achieved. The angular resolution for starting tracks is shown in Figure 2 (dashed) and is ∼1° in the interesting energy region; the reconstruction is worse than for through-going events (solid), due to a smaller lever arm for tracks starting within the fiducial volume of the detector.
Citation: Gaisser et al. 2014
Citation function: Background
NLFFF extrapolation provides the reconstructed coronal magnetic field for AR 11158 from 2011 February 13 − 2011 February 15 (Figures 1(d)–(f)). The field lines (yellow lines) within the core of the AR have arcade-like structure with a relatively strong twist mainly near the PIL. These figures show that the magnetic field evolved during this period. Although we did not quantitatively compare the field lines with the observation, in general, the reconstructed coronal field morphologies match with the observations in Figures 1(a)–(c). The general morphologies and the locations of the high-twist fields are also in agreement with many previous studies (Jing et al. 2012; Sun et al. 2012; Dalmasse et al. 2013; Inoue et al. 2013, 2014a; Liu et al. 2013; Wang et al. 2013; Aschwanden et al. 2014; Malanushenko et al. 2014; Zhao et al. 2014). Unlike Zhao et al. (2014) who could identify the twisted flux rope from the topology of the reconstructed coronal field, we could not find an obvious topological signature of a flux rope existing in our NLFFF during our analysis time window. It might be due to the fact that there was little magnetic flux with twist higher than one turn in our NLFFF and it is difficult to topologically define it as a flux rope. However, our result is consistent with other NLFFF results (Jing et al. 2012; Sun et al. 2012; Liu et al. 2013; Wang et al. 2013; Inoue et al. 2014a; Malanushenko et al. 2014). The high-twist region in our result is also in agreement with the region with high helicity flux (Dalmasse et al. 2013) and the location of the flare ribbons (Bamba et al. 2013; Liu et al. 2013), as well as the high current density region (Janvier et al. 2014). Figures 1(g)–(i) show the evolution of the twist distribution map, with the magnetic twist of the field lines plotted at the footpoints of field lines according to a color scale. This shows that the high-twist (strongly right-handed twist corresponding to Tw > 0.5) areas are concentrated in only a limited part of the AR. The high-twist area grew and became even more twisted just before the X2.2 flare (Figure 1(i)). Most parts of the AR have twist values less than 0.25, but near the PIL the twist can reach more than 0.5, even up to about a full turn. This is consistent with the results of Sun et al. (2012) and Inoue et al. (2014a). A high-twist (strong negative/left-handed twist) area also developed in the eastern part of the AR, which did not exist initially on February 13. Both of these high-twist areas produced several flares. However, here we focus on the flares that resulted from the high-twist core region near the center of the AR, where the M6.6 and X2.2 flares occurred.
Citation: Liu et al. 2013
Citation function: Similarities
Citation: Liu et al. 2013
Citation function: Similarities; Similarities
Citation: Liu et al. 2013
Citation function: Similarities
Citation: Liu et al. 2013
Citation function: Similarities
Citation: Liu et al. 2013
Citation function: Similarities; Similarities
Citation: Liu et al. 2013
Citation function: Similarities
A number of semi-analytic models (SAMs) have attempted to reproduce submm number counts (e.g. Granato et al. 2000; Fontanot et al. 2007; Somerville et al. 2012). One such model is the galform (SAM), which has been tuned to successfully reproduce the number counts of 850 μm and $\mathrm{1.1 \, mm}$ selected galaxies.1 However, in order to achieve this good agreement galform invokes a top-heavy initial mass function (IMF). Early versions of the model used a flat IMF above $1 \, \mathrm{M_{\odot }}$, in sub-L* mergers (Baugh et al. 2005; Swinbank et al. 2008). This is required to produce sufficiently bright submm emission during frequent low-mass merger events. Later versions of the model used a more moderately top-heavy IMF in starbursts, triggered by disc instabilities rather than mergers, and found similarly good agreement with the number counts (Cowley et al. 2015, 2019; Lacey et al. 2016; Park et al. 2016). However, such IMF variability is still controversial, particularly extreme forms and any dependence on merger state (Bastian, Covey & Meyer 2010; Hopkins 2013; Krumholz 2014), and is inconsistent with the constraints on the IMF in massive star-forming galaxies that are significantly less extreme (e.g. Tacconi et al. 2008), though there is tentative evidence of a bottom-light/top-heavy IMF in both local star-forming region analogues (Motte et al. 2018; Schneider et al. 2018) and some gravitationally lensed high-redshift starbursts (Zhang et al. 2018). Safarzadeh, Lu & Hayward (2017) showed that a variable IMF is degenerate with a number of other modelling processes in SAMs, such as the form of stellar feedback. They highlight that taking in to account dust mass allows for a good fit to the number counts without resorting to a variable IMF. Most recently, the shark SAM (Lagos et al. 2018) is able to broadly reproduce the 850 μm counts (whilst slightly overestimating the bright-end counts compared to S2CLS; Geach et al. 2017) using a fixed Chabrier (2003) IMF (Lagos et al. 2019). They attribute the good agreement to their use of physically motivated attenuation curves obtained from a self-consistent galaxy evolution model (eagle; Trayford et al. 2020).
Citation: Cowley et al. 2015
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The pulsational analysis presented in this work makes use of full stellar evolution models of pre-WDs generated with the LPCODE stellar evolution code. LPCODE computes in detail the complete evolutionary stages leading to WD formation, allowing the WD and pre-WD evolution to be studied in a consistent way based on the evolutionary history of progenitors. Details of LPCODE can be found in Althaus et al. (2005, 2009, 2013) and references therein. Here, we mention only those ingredients employed which are relevant for our analysis of low-mass, He-core WD and pre-WD stars (see Althaus et al. 2013, for details). The standard Mixing Length Theory (MLT) for convection in the version ML2 is used (Tassoul et al. 1990). In this prescription, due to Bohm & Cassinelli (1971), the parameter α (the mixing length in units of the local pressure scale height) is set equal to 1, while the coefficients a,b,c that appear in the equations for the average speed of the convective cell, the average convective flux, and the convective efficiency (see Cox 1968), have values a = 1,b = 2,c = 16. We emphasize that the results presented in this work are insensitive to the prescription of the MLT employed. In particular, we have also used the ML1 (α = 1,a = 1 / 8,b = 1 / 2,c = 24, Böhm-Vitense 1958) and ML3 (α = 2,a = 1,b = 2,c = 16, Tassoul et al. 1990) recipes, and we obtain the same results than for ML2. The metallicity of the progenitor stars has been assumed to be Z = 0.01. It is worth mentioning that the pulsation stability results presented in this paper do not depend on the value of Z2. Radiative opacities for arbitrary metallicity in the range from 0 to 0.1 are from the OPAL project (Iglesias & Rogers 1996). Conductive opacities are those of Cassisi et al. (2007). The equation of state during the main sequence evolution is that of OPAL for H- and He-rich compositions. Neutrino emission rates for pair, photo, and bremsstrahlung processes have been taken from Itoh et al. (1996), and for plasma processes we included the treatment of Haft et al. (1994). For the WD regime we have employed an updated version of the Magni & Mazzitelli (1979) equation of state. The nuclear network takes into account 16 elements and 34 thermonuclear reaction rates for pp-chains, CNO bi-cycle, He burning, and C ignition. Time-dependent diffusion due to gravitational settling and chemical and thermal diffusion of nuclear species has been taken into account following the multicomponent gas treatment of Burgers (1969). Abundance changes have been computed according to element diffusion, nuclear reactions, and convective mixing. This detailed treatment of abundance changes by different processes during the WD regime constitutes a key aspect in the evaluation of the importance of residual nuclear burning for the cooling of low-mass WDs.
Citation: Böhm-Vitense 1958
Citation function: Uses
We used the stellar population synthesis code STARLIGHT (Cid Fernandes et al. 2004) to describe the age distributions and metallicities of the stellar populations that fit the integrated light spectrum of object I. For object H it is not possible to perform the stellar population synthesis because of the low signal-to-noise ratio of the spectrum. This code is extensively discussed in Cid Fernandes et al. (2004; 2005), and it is built upon computational techniques that originally were developed for empirical population synthesis with additional ingredients from evolutionary synthesis models. This method was used by Krabbe et al. (2011) and Faúndez-Abans et al. (2012, 2013) and has been successful in describing the stellar population in interacting galaxies. Briefly, the code fits an observed spectrum with a combination of N single stellar populations (SSPs) from the models of Bruzual & Charlot (2003). These models are based on a high-resolution library of observed stellar spectra, which allows for detailed spectral evolution of the SSPs across the wavelength range of 3200−9500 Å  with a wide range of metallicities. We used the Padova 1994 tracks, as recommended by Bruzual & Charlot (2003), with the Salpeter initial mass function (Salpeter 1955). Extinction is modeled by STARLIGHT as due to foreground dust, using the Large Magellanic Cloud average reddening law of Gordon et al. (2003) with RV = 3.1, and parametrized by the V-band extinction AV. The SSPs used in this work cover 15 ages, t = 0.001, 0.003, 0.005, 0.01, 0.025, 0.04, 0.1, 0.3, 0.6, 0.9, 1.4, 2.5, 5, 11, and 13 Gyr, and three metallicities, Z = 0.2Z⊙, 1 Z⊙, and 2.5 Z⊙, adding to 45 SSP components. The fitting is carried out using a simulated annealing plus Metropolis scheme, with regions around emission lines and bad pixels excluded from the analysis. The upper panel of Fig. 4 shows the observed spectrum corrected for reddening and the model stellar population integrated light spectrum of object I. The synthesized spectrum fits the observed one closely. The results indicate that this galaxy is an old object, with 100% of the flux contribution at λ5870 Å being provided by a stellar population of 13 Gyr, with Z = 2.5Z⊙ and AV = 0.36. This spectrum is dominated principally by the bulge signal and does not represent the whole galaxy; the age estimate then is valid only for its bulge.
Citation: Padova 1994
Citation function: Uses; Uses
The (sub)mm wavelength range is of particular interest for various reasons: at sufficiently long wavelengths, disks become optically thin, and an estimate of their dust mass can be directly obtained (via some assumptions) by simply measuring their flux (e.g., Beckwith et al. 1990). Although the bulk of the disk mass in the system is in gaseous phase, fiducial (or measured, when available) gas-to-dust ratios provide an indirect estimate of the total mass in the disk. This is a crucial parameter for planet formation theories because it determines the available reservoir for this process. Using this method, surveys of star-forming regions with (sub)mm facilities such as SMA and ALMA have determined that protoplanetary disks have typical masses of 0.1%–0.5% of that of their host star (e.g., Andrews & Williams 2005; Andrews et al. 2013; Pascucci et al. 2016). On the other hand, dust growth represents the initial stage of planet formation; the observed spectral index at these wavelengths can be linked to the dust opacity in the disk, Informative of its properties and grain sizes (e.g., Miyake & Nakagawa 1993; Draine 2006). In fact, the comparison of the millimeter spectral index of the interstellar medium (ISM) with that of protoplanetary disks has already revealed significant dust growth in these disks, implying the presence of mm/cm-sized grains in many of them (e.g., D’Alessio et al. 2001; Lommen et al. 2010; Ricci et al. 2010a, 2010b; Ubach et al. 2012). The combination of the mm spectral index with additional information at other wavelengths, such as the spectral index at near/mid infrared (IR) wavelengths or silicate features may also point to links between the evolution of the inner and outer regions of the disks. As an example, Lommen et al. (2010) identified a tentative correlation between the strength of the 10 μm silicate feature and the 1–3 mm spectral index for a sample of T Tauri and Herbig Ae/Be stars, suggesting a connection between the evolution of the inner and outer regions of disks, although a later study by Ricci et al. (2010b) found no signs of such a correlation for disks in the Taurus and Ophiuchus star-forming regions. Despite the obvious interest of this wavelength regime, disks have relatively weak emission at millimeter wavelengths and many of them currently lack this type of data (or, at least, sufficient observations to provide robust estimates of their spectral indices).
Citation: Pascucci et al. 2016
Citation function: Background
To minimize the undesired effects caused by observing different lunar locations, we conducted observations according to the following procedure. On each observing night, we first pointed the Nayuta telescope toward the crater Grimaldi (selenographic coordinate: 68.6°W, 5.2°S) in the waxing phase and the crater Neper (84.5°E, 8.8°N, east of Mare Crisium) in the waning phase, after correcting the pointing error measured using a nearby star. Both craters are near the lunar edge (distances ≲2′). Then, we scanned the Moon along the RA axis until the edge of the Moon was placed near the center of the field of view (FOV). An example of the observed (and reduced) images is shown in Fig. 2. Our target locations are not on a major maria and near sites repeatedly observed in previous Earthshine photometry because they were expected to have roughly comparable albedos (Qiu et al. 2003; Pallé et al. 2004; Montañés-Rodríguez et al. 2007). Half of the FOV was reserved for the sky, which allows the sky background intensities and their positional gradients to be measured. The position angle of the instrument (ϕinspa) was maintained at 90° from the equatorial north, as measured counter-clockwise, so that the long side of the FOV was aligned with the RA axis. Telescope tracking was conducted in accordance with the sky motion of the Moon, which was calculated at the Jet Propulsion Laboratory (JPL) Horizons system1. Because the tracking was not perfect, we shifted the telescope east or west with a typical interval of ~ 30 min so that the lunar edge remained near the center of the FOV. Features on the Moon were hardly recognizable in the raw images because of the dim Earthshine and strong scattered light from the day side of the Moon, though we were able to visually identify the lunar edge in most cases2. Despite our efforts, the actually observed location may have varied night by night even within one phase (waxing phase or waning phase), or on an hourly timescale during a single night. Possible impacts induced by different lunar locations (namely different degrees of depolarization) are discussed in Sect. 4.1.
Citation: Pallé et al. 2004
Citation function: Motivation
Currently, about 20 Be stars are known to have stripped companions, most of which were confirmed with far-UV (FUV) spectroscopy, as it is in the FUV where the flux ratios are most favorable (Wang et al. 2021; Klement et al. 2022, and references therein). All of the spectroscopic FUV detections were found to be compatible with the sdO nature of the companions, while no firm case of a cooler sdB companion has been presented as of yet. There are also more than 160 confirmed and candidate Be X-ray binaries (BeXRBs; Raguzova & Popov 2005),
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List updated at http://xray.sai.msu.ru/~raguzova/BeXcat/. which are mostly Be+NS systems that probably evolved in a similar fashion as the Be+sdO binaries but from more massive progenitor systems (Reig 2011). The Be primaries occupy a narrow range between O9 and B2 in spectral type (Reig et al. 2017). These systems are conspicuous due to the X-ray emission resulting from the (episodic) accretion of Be disk material onto the compact object, so that the sample is drawn from a much larger (partly extragalactic) volume. The WD companions to Be stars proved unexpectedly elusive (Meurs et al. 1992; Cohen et al. 1997), but several supersoft X-ray emission sources consistent with (early-)Be+WD systems undergoing a type II BeXRB outburst were recently detected in the Magellanic Clouds (Coe et al. 2020; Kennea et al. 2021). One Be star (MWC 656) was reported to have a BH companion, but this was recently shown to be questionable on the basis of new higher-quality spectra, which rather point toward another Be+sdO system (Rivinius et al. 2022). Object HD 93521 is the first candidate postmerger Be star (Gies et al. 2022), and members of a Be star runaway population were found using Hipparcos and Gaia astrometric catalogs (Berger & Gies 2001; Boubert & Evans 2018; Wang et al. 2022). On the other hand, Be stars have been found missing among B stars at high Galactic latitude (Martin2004, 2006). Meanwhile, not a single early-type Be star has been confirmed to have a close MS companion (Gies 2000; Bodensteiner et al. 2020). However, the recently studied case of the B6Ve star α Eri, which is a highly eccentric binary with an early A-type dwarf companion on a 7 yr orbit, appears to be the first confirmed case of a Be star that does not owe its nature to mass transfer in a close binary, as the presence of a close stripped companion was ruled out (Kervella et al. 2022b). This implies that two evolutionary channels—single and binary—indeed exist for the formation of Be stars.
Citation: Cohen et al. 1997
Citation function: Background
During a GF, a huge amount of the magnetic energy E > 1046 erg is subsequently released at the surface of the magnetar within the first 1 s, leading to the formation of a hot fireball similar to the case of classic GRBs (Mészáros & Rees 2000), which might be rich in electron–positron pairs (Thompson & Duncan 1995; Fermi-LAT Collaboration 2021). The afterglow emission of a GRB is generally well explained by the synchrotron emission from electrons accelerated by the shock produced during a relativistic ejecta colliding with an external medium. The ejecta might be formed by the baryon-rich crust of the NSs undergoing a large-scale disruption, or a dense outflowing gas of electron–positron pairs and radiations (Thompson 2021). Zhang et al. (2020) prefer a high fraction of electron energy for GRB 200415A, which is consistent with the situation of an electron–positron pair-dominating ejecta. Chand et al. (2021) suggest that a baryon-poor outflow can explain the high-energy afterglow emission of GRB 200415A, while a baryon-rich outflow is also viable if the dissipation happens below the photosphere via internal shocks. Therefore, in this paper, we calculate two situations that the ejecta is dominated by electron–positron pairs or protons, respectively. Around the magnetar, a cavity might be created by the pulsar wind and the earlier SGR activity. Holcomb et al. (2014) suggest that the Poynting flux from a pulsar can evacuate a cavity with a size of sub-parsec. Such a cavity is also needed to explain the radio afterglow of the GF from SGR 1806-20 (Gaensler et al. 2005; Gelfand et al. 2005; Taylor et al. 2005; Granot et al. 2006). Fermi-LAT Collaboration (2021) suggests a cavity environment to explain the expansion of the ejecta before its colliding with the bow-shock shell for GRB 200415A. Calculations in Zhang et al. (2020) on the afterglow of GRB 200415A also suggest a low density of ambient medium, which is consistent with the cavity environment. In this paper, we discuss two situations for the CSM density, including a density of $n = 10^{-2}\ \rm cm^{-3}$, and a low density of $n = 10^{-5}\ \rm cm^{-3}$ for the situation of a cavity environment surrounding the magnetar.
Citation: Chand et al. (2021)
Citation function: Motivation