adsabs/FOCAL huggingface dataset:
The λ Boötis stars are a group of chemically peculiar objects on the upper main-sequence, showing underabundances (~1–2 dex) of iron-peak elements and near-solar abundances of C, N, O, and S (e.g., Kamp et al. 2001; Heiter 2002; Andrievsky et al. 2002). The class was discovered by Morgan et al. (1943) and named following the bright prototype λ Boötis, which is one extreme member of the class. This group of refractory-poor objects comprises about 2% of early B through early F stars (Gray & Corbally 1998; Paunzen 2001). However, among the pre-main-sequence Herbig Ae/Be stars, that is to say the progenitors of A-type stars, the λ Boötis-like fraction is about 33% (Folsom et al. 2012). The origin of the peculiarity still remains a puzzle, as we can read in the recent discussion of Murphy & Paunzen (2017). Unlike common chemical peculiarities seen in Am and Ap stars, λ Boötis stars are not constrained to slow rotation (Abt & Morrell 1995; Murphy et al. 2015). Cowley et al. (1982) first suggested that λ Boötis could possibly originate from the interstellar medium (ISM) with a nonsolar composition, or from the separation of grains and gas. Then, Venn & Lambert (1990) proposed that λ Boötis stars likely occur when circumstellar gas is separated from grains and then accreted to the stars due to the similarity ofits abundance pattern with the ISM. Other proposed mechanisms include the interaction of a star with a diffuse interstellar cloud (Kamp & Paunzen 2002; Martinez-Galarza et al. 2009), where the underabundances are produced by different amounts of accreted material. Turcotte & Charbonneau (1993) estimated that once the accretion stops, the photospheric mixing and meridional circulation would erase this peculiar signature on a ~1 Myr timescale. By studying thedistribution of λ Boötis stars on the HR diagram, Murphy & Paunzen (2017) conclude that multiple mechanisms could result in a λ Boötis spectra, depending on the age and environment of the star.
Citation: Murphy et al. 2015
Citation function: Compare/Contrast
Feedback from massive stars plays a critical role in the star formation process and evolution of molecular clouds. In particular, expanding H ii regions may have a positive effect on star formation, i.e., they can trigger a new generation of star formation in molecular clouds either by sweeping ambient clouds into dense shells or by compressing nearby dense clouds into bound clumps/cores (for details, see Deharveng et al. 2010). In both cases, dense material eventually fragments to form new stars. Since the PG108.3 cloud consists of three evolved H ii regions, one can thus speculate that the formation of young sources in the region might have been triggered by the expanding bubble. However, our observations largely do not favor such a process because in such a scenario one would expect the distribution of young YSOs or star-forming cores at the outskirts of the H ii regions (see, e.g., Zavagno et al. 2006; Jose et al. 2013; Panwar et al. 2014; Liu et al. 2016; Samal et al. 2018). In contrast, we find that no young Class I sources around S147 and most of Class I sources around S148 are cospatial with the Class II sources located near the center of the ionized gas. Since young clusters often harbor Class III to Class I sources as a part of the cluster formation process (e.g., Jose et al. 2017; Panwar et al. 2017, 2018), the Class I YSOs of the S148 are thus most likely part of the central cluster responsible for the ionization of S148. Given the fact that the H ii region S149 is located near the boundary of S148 and smaller in size, one can argue that S149 might have been influenced by S148. Unfortunately, from the current data we have no way to estimate precisely either the age of the H ii regions or the embedded YSOs in S148. Nonetheless, we searched for signatures of the early phases of high-mass star formation such as masers and outflows in the vicinity of S148, and our search resulted in no such observations. The small size of the S149 H ii regions could be due to the fact that it is ionized by a less massive star compared to S148, and as a result, S149 is expanding at a slow rate compared to S148, assuming that both of them formed in a similar environment. Considering the fact that S149 is optically visible, its dynamical age is close to the age of S148, and the ionized gas pressures of both regions are within a factor of two, we regard the probability that the formation of S149 is solely due to expansion of S148 to be rather low, although we cannot ignore the possibility that the expanding ionization front of S148 might have accelerated star formation in S149 after the initial clump formation and fragmentation. A way to prove such a hypothesis is to compare the molecular gas pressure of S149 with external gas pressure generated by S148 (e.g., Thompson et al. 2004; Kim et al. 2017; Liu et al. 2017) in the early stage of star formation; however, in S149, neither do we find significant cold gas nor do Azimlu & Fich (2011) find any molecular clumps to obtain some clue in this direction.
Citation: Liu et al. 2016
Citation function: Compare/Contrast; Compare/Contrast
CCSNe are also of interest for GW astronomy as targets in their own right. As the sensitivity of GW detectors increases, they will begin to detect not only binary mergers but also other lower amplitude sources of GWs such as CCSNe. Accurate knowledge of the GW emission from CCSNe will be essential for detection and parameter estimation. The GW signal from rotational core bounce has already been well covered in the literature (e.g. Dimmelmeier et al. 2008; Abdikamalov et al. 2014; Fuller et al. 2015; Richers et al. 2017). In the non-rotating case, the GW emission from the post-bounce phase has been studied using self-consistent 3D simulations by many groups (Kuroda, Kotake & Takiwaki 2016; Andresen et al. 2017, 2019; Kuroda et al. 2017, 2018; Powell & Müller 2019, 2020; Radice et al. 2019; Andresen, Glas & Janka 2020; Mezzacappa et al. 2020; Pan et al. 2020). The structure of the GW emission has shown common features in different simulations from recent years. The dominant emission feature in the GW emission is due to the quadrupolar surface f/g mode 1 of the proto-neutron star (PNS), which produces GW frequencies rising in time from a few hundred Hz up to a few kHz (Müller, Janka & Wongwathanarat 2012; Sotani et al. 2017; Kuroda et al. 2018; Morozova et al. 2018; Torres-Forné et al. 2018, 2019). In addition, some models (Kuroda et al. 2016, 2017; Andresen et al. 2017; Mezzacappa et al. 2020; Powell & Müller 2020) exhibit low-frequency GW emission due to the standing accretion shock instability (SASI; Blondin, Mezzacappa & DeMarino 2003; Blondin & Mezzacappa 2006; Foglizzo et al. 2007). In rapidly rotating models, very strong GW emission can also occur during the post-bounce phase due to a corotation instability (Takiwaki & Kotake 2018). The emerging understanding of the GW emission features has led to the formulation of universal relations for the GW emission (Torres-Forné et al. 2019) and paved the way for phenomenological modelling for CCSN signals (Astone et al. 2018). Further work is still needed, however, to extend these models to fully explore CCSN GW signals from across the progenitor parameter space. The majority of 3D simulations that include GW emission are for progenitor stars below $30\, \mathrm{M}_{\odot }$. In this paper, we perform simulations of high-mass Population III (Pop-III) stars in the pulsational pair instability regime to expand the parameter space coverage of 3D simulations and to provide further insights into the massive and very massive star remnant BH population.
Citation: Andresen et al. 2017
Citation function: Background
Citation: Andresen et al. 2017
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: Illarionov & Sunyaev 1975
Citation function: Uses; Uses
Solar vortex tubes can be spontaneously generated by turbulent convection. In simulations of quiet Sun regions, vortices are found along intergranular lanes (Shelyag et al. 2011a; Kitiashvili et al. 2012; Moll et al. 2012; Silva et al. 2020). These structures have an average lifetime of around 80 s (Silva et al. 2021) and a radius between 40 and 80 km (Shelyag et al. 2013; Silva et al. 2020). Solar kinetic vortex tubes (Silva et al. 2021) act as a sink for magnetic field, creating magnetic flux tubes that expand with height (Kitiashvili et al. 2012; Moll et al. 2012; Silva et al. 2020). The concentration of magnetic flux leads to a high magnetic field tension, which can prevent the magnetic field lines from being twisted by the rotational motion (Shelyag et al. 2011b; Moll et al. 2012; Nelson et al. 2013; Silva et al. 2021). In some cases, twisted magnetic flux tubes appear close enough to flow vortices, leading to magnetic and kinetic vortex structures closely coexisting in regions with high plasma-β (Wedemeyer & Steiner 2014; Rappazzo et al. 2019; Silva et al. 2021). The vortical motions can still trigger perturbations along magnetic lines that could lead to wave excitation, e.g., Battaglia et al. (2021). The vorticity evolution in the magnetized solar atmosphere is mainly ruled by the magnetic field, which also influences the general shape of vortices (Shelyag et al. 2011a). Based on the analysis of swirling strength, the part of the vorticity only linked to swirling motion (Shelyag et al. 2011b; Canivete Cuissa & Steiner 2020) showed that the magnetic terms in the swirling equation evolution tend to cancel the hydrodynamic terms close to the solar surface, whereas the magnetic terms dominate alone the production of swirling motion in the chromosphere. The magnetic field also tends to play an important role in the plasma dynamics along the whole vortex tube, as the Lorentz force has a magnitude comparable to the pressure gradient (Silva et al. 2020; Kitiashvili et al. 2013). High-speed flow jets have also been linked to simulated vortex tubes, driven by high-pressure gradients close to the photosphere and by Lorentz force in the weakly magnetized upper solar photosphere (Kitiashvili et al. 2013). In general, the averaged radial profile of magnetic field, angular velocity, pressure gradient inside of the vortex tube at the lower chromosphere and photosphere levels show similar behavior (Silva et al. 2020).
Citation: Wedemeyer & Steiner 2014
Citation function: Background
Citation: Rappazzo et al. 2019
Citation function: Background
Since the discovery of the radio burst, there have been extensive follow-up observations of SGR J1935+2154 across the electromagnetic spectrum. The lack of another radio pulse coincident with an X-ray flare puts interesting constraints on the emission mechanism and begs the question of whether we should be able to see such radio bursts in other active Galactic magnetars. In this context, connecting FRBs with extragalactic magnetars is tantalizing. Current theories that propose FRB emission from a magnetar can be broadly divided into two categories: (1) far-away models, where the FRB is generated by a maser away from the neutron star, and (2) close-in models where the FRB is produced in the magnetospehere of the star. The maser emission model runs into difficulties when explaining all the observed radio and X-ray properties of the contemporaneous radio/X-ray burst seen from SGR J1935+2154 (see Lu, Kumar & Zhang 2020, for more details). Younes et al. (2020a) have shown that the X-ray burst contemporaneous with the radio burst was spectrally unique compared to all other burst in the activity period and it also supports it having a polar cap origin. Hence, if we assume that FRBs produced by magnetars are created in the magnetosphere close to the polar cap, we can expect them to be significantly beamed (Lu, Kumar & Zhang 2020). This of course also means that the source must still be exhibiting bursts infrequently in the radio and that more of them might be associated with an X-ray burst than we observe. Observational evidence so far does suggest a connection between the X-ray and the radio emission mechanisms prevalent in neutron stars and any changes in one of them affects the other (Archibald et al. 2017). It is believed that while the X-ray bursts and radio pulsations in these sources come from different regions in the magnetosphere, the pair plasma causing the X-ray flares can affect the acceleration of radio-emitting particles. This was shown in PSR J1119−6127, a high B-field radio pulsar where a series of X-ray bursts from the source quenched the radio pulsations with the radio emission returning a few minutes after the last X-ray burst (Archibald et al. 2017).
Citation: Lu, Kumar & Zhang 2020
Citation function: Compare/Contrast
Einstein’s theory has been tested successfully mainly in the regime of weak gravity through Solar system tests and laboratory experiments. But the validity of this highly accepted theory still faces stringent constraint in the regime of strong gravity, viz., the region near to a black hole, ultradense compact stars, and expanding universe. The recent discovery of peculiar highly overluminous SNeIa, e.g. SN 2003fg, SN 2006gz, SN 2007if, and SN 2009dc (Howel et al. 2006; Scalzo et al. 2010) indicates a huge Ni-mass and confirms the highly super-Chandrasekhar white dwarfs, having mass 2.1–2.8 M⊙, as a suitable progenitors (Howel et al. 2006; Hicken et al. 2007; Yamanaka et al. 2009; Scalzo et al. 2010; Silverman et al. 2011; Taubenberger et al. 2011). Recently, Linares, Shahbaz & Casares (2018) have discovered a highly massive pulsar of mass $2.27^{+0.17}_{-0.15}~\mathrm{M}_{\odot }$ in their observation of compact binaries PSR J2215+5135. Clearly, these observations are not only questioning the standard Chandrasekhar limit for the compact stellar objects but also invoking the necessity of modification of GR in the strong gravity regime. Interestingly, our study reveals that due to the effect of $f\left(R,\mathcal {T}\right)$ gravity theory, the maximal mass limits rise higher than their standard values in GR for the chosen parametric values of χ. Hence, the stellar systems in the framework of $f\left(R,\mathcal {T}\right)$ gravity theory may also explain the observed massive stellar systems, viz., massive pulsars, super-Chandrasekhar stars, and magnetars, etc., which GR hardly can explain suitably so far. In support of the achieved result in the present investigation, it is worth mentioning that the important study by De Laurentis (2018) also reveals that the application of the Noether Symmetry Approach can explain the extreme massive stars which are supposed to be unstable in the framework of GR. However, in the limit χ = 0, one may retrieve the solutions of the standard Einstein gravity.
Citation: Yamanaka et al. 2009
Citation function: Background; Background
Solar cycle predictions are needed to plan long-term space missions and are of high importance for space weather applications. Currently, precursor methods are the most favored models for the prediction of solar cycle strength (Conway 1998; Svalgaard et al. 2005; Kane 2008; Hathaway 2009). These precursor techniques often use geomagnetic activity levels near or before the time of solar cycle minimum (Ohl & Ohl 1979; Feynman 1982; Gonzalez & Schatten 1988; Thompson 1993; Wilson et al. 1998). Predicting the amplitude of a solar cycle can be done using solar polar magnetic fields from the previous cycle as “precursors” of the next cycle (Schatten & Sofia 1987; Schatten 2005; Svalgaard et al. 2005; Wang & Sheeley 2009; Muñoz-Jaramillo et al. 2012). The other class of precursor techniques that do not need an a priori physical understanding of the causal relations (i.e., that do not require any knowledge of the physics involved) is based on finding particular sunspot number characteristics that serve as indicators of the next cycle strength (Ramaswamy 1977; Lantos 2006; Cameron & Schüssler 2008; Brajša et al. 2009). A number of techniques are used to predict the amplitude of a cycle during the time near and before sunspot minimum (Hathaway 2010). The two precursor types that have received the most attention are polar field precursors and geomagnetic precursors. The strength of the solar polar magnetic fields at solar minimum is a very accurate indicator of the maximum amplitude of the following solar cycle. Forecasts using the polar field method have proven to be consistently in the right range for cycles 21, 22, and 23 (Schatten & Sofia 1987; Schatten et al. 1996). The polar fields reach their maximal amplitude near the minima of the sunspot cycle. However, the maxima of the polar field curves are often rather flat, so approximate forecasts are feasible several years before the actual minimum (Hathaway 2010; Petrovay 2010). Using the rather flat and low maximum in polar field strength, Svalgaard et al. (2005) have been able to predict a relatively weak cycle 24. Such an early prediction is not always possible: early polar field predictions of cycles 22 and 23 had to be corrected later and only forecasts made shortly before the actual minimum finally converged (Hathaway 2010; Petrovay 2010).
Citation: Muñoz-Jaramillo et al. 2012
Citation function: Background
To confirm the importance of VO, water and an inversion layer (Evans et al. 2018; Mikal-Evans et al. 2019, 2020) obtained repeated HST observations of the transmission spectrum and the secondary eclipse using the STIS and WFC 3 instruments. The optical transmission spectrum displays rich variation, with multiple features consistent with VO absorption that Evans et al. (2018) could reproduce by assuming an isothermal T-P profile at 1500 K and a metallicity equivalent to 10× to 30× solar. Absorption bands of TiO appeared to be muted in the transmission spectrum, which was explained by Evans et al. (2018) as evidence of condensation of Ti-bearing species, which commences at higher temperatures than condensation of V-bearing species, producing for example, calcium titanates (Lodders 2002) while VO remains in the gas phase. Mikal-Evans et al. (2019) observed the day side emission spectrum with the G102 grism of WFC3 (0.8–1.1 μm), augmenting their earlier observations with the G141 grism. The G102 spectrum does not show the VO bands expected to be present there, and this led Mikal-Evans et al. (2019) to question the interpretation that the 1.2 μm feature is caused by VO emission. The secondary eclipse was observed at 2 μm (Kovács & Kovács 2019) and at optical wavelengths with the TESS instrument. These were analysed together with the preceding Hubble, Spitzer, and ground-based observations to yield tighter constraints on the atmospheric structure, composition and overall system parameters (Bourrier et al. 2020a; Daylan et al. 2019). These studies found that the hottest point on the day side exceeds a temperature of 3000 K, that the atmosphere is inverted on the day side, and a metallicity that is consistent with solar (Bourrier et al. 2020a) or slightly elevated (Daylan et al. 2019). Although the chemical retrievals follow different strategies (equilibriumversus free-chemistry), both indicate that a depletion of TiO relative to VO is needed to explain the observed emission spectrum, supporting the earlier findings by Mikal-Evans et al. (2019). Recently, Mikal-Evans et al. (2020) obtained new secondary-eclipse observations using the G141 grism of WFC3. Although confirming the presence of emission by H2O, a joint analysis with their previous WFC3 observations did not reproduce the emission feature at 1.2 μm, prompting the authors to entirely discard their previous interpretation of emission caused by VO.
Citation: Mikal-Evans et al. 2019
Citation function: Background; Background; Compare/Contrast
Citation: Mikal-Evans et al. (2019)
Citation function: Similarities
Galaxies seem to retain a memory of their spin orientation with respect to the cosmic web filaments and walls, as suggested by the results from large-scale cosmological hydrodynamical simulations (Dubois et al. 2014; Codis et al. 2018; Wang et al. 2018; Ganeshaiah Veena et al. 2019; Kraljic, Davé & Pichon 2020b). The mass dependence of the spin alignment signal is however debated. While some works confirmed the existence of a galaxy spin transition from parallel to perpendicular with respect to the filament’s direction (Dubois et al. 2014; Codis et al. 2018; Kraljic et al. 2020b), and analogously with respect to walls (Codis et al. 2018; Kraljic et al. 2020b), others (Ganeshaiah Veena et al. 2019; Krolewski et al. 2019) found preferential perpendicular orientation with respect to filaments at all masses with no sign of a spin transition. A possible interpretation of this lack of detection of a clear transition is the nature of the filaments, with galaxies in thinner filaments having their spins more likely perpendicular to the filament’s axis, compared to galaxies of similar mass in thicker filaments (Ganeshaiah Veena et al. 2019). This can be in turn understood recalling the multiscale nature of the problem and the conditional TTT (Codis et al. 2015) predicting larger transition mass for denser, thus thicker, filaments. Further support for this interpretation was provided by the findings of stronger impact of large-scale tides on the galaxy spin orientation in denser filaments (Kraljic et al. 2020b, using filament density as a proxy for the thickness of filaments). In addition to the stellar mass, the spin-filament alignment was shown to depend on other internal properties of galaxies. Blue or rotation-supported galaxies were found to dominate the alignment signal at low stellar mass, while red or dispersion-dominated galaxies tend to show a preferential perpendicular alignment (Codis et al. 2018; Wang et al. 2018; Kraljic et al. 2020b).
Citation: Kraljic et al. 2020b
Citation function: Background
Citation: Kraljic et al. 2020b
Citation function: Similarities; Background