adsabs/FOCAL huggingface dataset:
The evolution of galaxies sensitively depends on the stellar initial mass function (IMF). The understanding of the IMF has changed rapidly in the past decade. Despite a direct conflict (Kroupa et al. 2013) with the theoretical expectation that the IMF should vary as the star-forming environment alters (e.g. ambient gas temperature, metallicity, density, and pressure dependence, as argued by Adams & Fatuzzo 1996; Larson 1998; Elmegreen 2004; Dib et al. 2007; Papadopoulos 2010) the observed IMFs of star clusters in the local Universe are consistent with no variation (Kroupa 2001, 2002; Bastian et al. 2010; Offner et al. 2014; Hopkins 2018). Thus, the universal and invariant canonical IMF assumption was widely applied. The more recent observations, which are able to probe physical regimes further from the solar and Galactic neighbourhood, have been consistently suggesting a variation of the galaxy-wide IMF (gwIMF1). For the distribution of low-mass stars, a bottom-heavy IMF (excess in the number of low-mass stars) in the inner regions of massive metal-rich elliptical galaxies is indicated by the galaxy mass-to-light ratio (Li et al. 2017), spectral analysis of stellar-mass sensitive features (Vazdekis et al. 1997, 2003; Cenarro et al. 2003; van Dokkum & Conroy 2010; Conroy & van Dokkum 2012; Ferreras et al. 2013; Martín-Navarro et al. 2015; La Barbera et al. 2017; Parikh et al. 2018) and lensing studies (Auger et al. 2010; Oldham & Auger 2018)2. For the distribution of massive stars, independent evidence strongly indicates a systematically varying IMF (galaxy photometry: Hoversten & Glazebrook 2008; Meurer et al. 2009; Lee et al. 2009; Gunawardhana et al. 2011, metal abundance of galaxy clusters: Renzini & Andreon 2014; Urban et al. 2017, isotope abundance: Romano et al. 2017; Zhang et al. 2018), being top-heavy (more massive stars than predicted when assuming the canonical IMF) when the SFR is high and/or when the metallicity is low, as summarized by Kroupa et al. (2013), Yan et al. (2017), and Jeřábková et al. (2018). We note that the IMF can be bottom-heavy and top-heavy at the same time (see Jeřábková et al. 2018).
Citation: Hopkins 2018
Citation function: Compare/Contrast; Compare/Contrast; Compare/Contrast
As mentioned in Section 1, once the sample of matching halos between the twin N-body simulations is established, we can study the halo−LSS correlations. We employ the Hessian matrix method used in many previous articles (e.g., Hahn et al. 2007a, 2007b; Zhang et al. 2009; Kang & Wang 2015) to define the matrix as:
4






H


ij


=





∂


2




ρ


s




x




∂


x


j


∂


x


i





.


This method is based on the smoothed density field 






ρ


s




x




 at the halo position (based on a more accurate and improved algorithm by Wang et al. 2020), which can be given by the Cloud-in-Cell (CIC) technique (MacNeice 1995). The smoothing length Rs is the only parameter of the CIC, which can be regarded as the typical scale of a halo LSS environment identified by the Hessian matrix method. Many previous works (e.g., Aragón-Calvo et al. 2007; Hahn et al. 2007a; Zhang et al. 2009; Codis et al. 2012; Hoffman et al. 2012; Libeskind et al. 2013; Trowland et al. 2013, but see also Libeskind et al. 2014) used a constant smoothing length. To determine which Rs value should be chosen, we test some LSS properties (e.g., the environment and eigenvectors) of matching halos for three fixed Rs: 2.5,  5, and  10 h−1 Mpc. We find that the halo−LSS correlation, as well as other main conclusions we make below, become stronger as Rs decreases, which is a reasonable result according to our previous discussions. Consequently Rs = 2.5 h−1 Mpc is chosen hereafter. The three eigenvalues of the Hessian matrix are marked as λ1, λ2, and λ3, with corresponding eigenvectors e1, e2, and e3. Eigenvalues λi can be used to define the LSS environment of dark matter halos according to the number of positive eigenvalues (Zel’Dovich 1970; Hahn et al. 2007a, 2007b; Libeskind et al. 2018; Zhang & Yang 2019); i.e.,

1.
void: 0  λ1 λ2 λ3


2.
wall: λ1  0  λ2 λ3


3.
filament: λ1 λ2  0  λ3


4.
cluster: λ1 λ2 λ3  0

and the eigenvectors ei stand for the three compressed directions of the smoothed density field. The e3 vector indicates the least compressed direction, which is a robust and universal definition of the LSS. In this work, we will focus on the alignment of halo spin and shape with the e3 vector; i.e., 




cos


θ


3


=
a
·


e


3




, where a is the halo spin or major-axis vector.
Citation: Aragón-Calvo et al. 2007
Citation function: Uses; Uses
The energy levels calculated by Fritzsche, Fischer & Fricke (1998) with grasp92 code are more closer to the NIST values as compared with other calculations within 3.6 per cent for the levels in the ground configurations. Our calculated MCDF energies agree well with NIST data within 6.8 per cent and the maximum error appears for ground state configuration (3s23p3) levels. The average percentage difference between our grasp1 energies and the measured values is 1.9 per cent. For the lowest 41 energy levels, we see that our calculated energies and previous fac energies (Landi & Bhatia 2010; Wang et al. 2012) belonging to the n = 3 complex in Ni xiv agree well with each other and both set of calculation are of the same accuracy. However, our calculated ground state configuration levels show better agreement with measured energies in comparison to those of fac calculations (7.2 per cent). When compared to MCDF calculations performed by Aggarrwal, Keenan & Msezane (2003) with grasp Dyall et al. (1989) code using a model of 14 configurations, our grasp1 levels in the excited configurations lie closer to the NIST ones. In the case of four energy levels (2D5/2,3/2,4P5/2,3/2), our calculated energy levels are more closer to NIST than those of Aggarrwal et al. (2003). The average percentage difference of these grasp energies with NIST values is 2.2 per cent, and the maximum error of 6.3 per cent appears for the first excited state (2D3/2) of ground configuration. For more accurate target states we adopt a model of 13 electronic configurations which yields target-states energies as closer to NIST values as possible. For instance, our calculated energy levels 2D3/2,5/2 associated with the configuration 3s3p4 lie more closer to the observed values as compared to existing calculations. The remaining set of energy levels are almost of the same accuracy as those of previous data when compared to experimental energies. From this comparison it is evident that present grasp1 calculated N-electron target wavefunctions demonstrate good agreement with experimental values, which is necessary for more accurate PI cross-sections of Ni xiii.
Citation: Wang et al. 2012
Citation function: Similarities; Similarities; Differences
Gamma Doradus (γ Dor) stars, with 1.4 M⊙ ≲ M* ≲ 2.0 M⊙, and slowly-pulsating B-type (SPB) stars, with 2.5 M⊙ ≲ M* ≲ 8 M⊙, exhibit high-order gravity-mode (g-mode) pulsations, gravito-inertial pulsations (Van Reeth et al. 2016), and/or purely inertial pulsations, such as r-modes (Saio et al. 2018). The restoring forces for gravity-modes and purely inertial pulsation modes are buoyancy and the Coriolis force, respectively. In the case of gravito-inertial pulsation modes, both forces contribute. As predicted by asymptotic theory, the pulsation periods for γ Dor and SPB stars were observed to form period spacing patterns (e.g. Chapellier et al. 2012; Chapellier & Mathias 2013; Kurtz et al. 2014; Bedding et al. 2015; Saio et al. 2015; Van Reeth et al. 2015; Ouazzani et al. 2017). The pulsation periods are equidistant in the asymptotic regime (with radial order n ≫ spherical degree l) for a non-rotating chemically homogeneous star (Tassoul 1980). Chemical gradients in the deep stellar interior cause pulsation mode trapping, which introduces non-uniform variations in the spacings (Miglio et al. 2008). On the other hand, the stellar rotation leads to shifts in the observed pulsation mode frequencies (Bouabid et al. 2013; Salmon et al. 2014; Van Reeth et al. 2015; Moravveji et al. 2016). For slowly rotating stars, the observed pulsation modes are split into frequency multiplets that depend on the mode identification. For moderate to fast rotators, that is, with rotation in the order of or more than 20% of the critical rotation rate, the observed period spacing patterns have a clear slope (e.g. Van Reeth et al. 2016; Ouazzani et al. 2017). Prograde (azimuthal order m >  0) and zonal modes have a downward slope, that is, the spacing between consecutive pulsation periods decreases with increasing pulsation period, i.e. radial order n. The period spacing patterns of retrograde modes (with m   0) mostly have an upward slope. For stars with detected period spacing patterns, this has been exploited to derive the near-core stellar rotation (e.g. Kurtz et al. 2014; Saio et al. 2015, 2018; Triana et al. 2015; Murphy et al. 2016; Schmid & Aerts 2016; Van Reeth et al. 2016; Ouazzani et al. 2017).
Citation: Van Reeth et al. 2015; Van Reeth et al. 2015
Citation function: Background; Motivation
With all this in mind, our last question is: What powers soft X-rays and [O III] in LINERs? This has no clear answer, but both are not tracing the same mechanism, since none of them match in morphologies. This is a clear difference between type-2 Seyferts and LINERs. In favour of the soft X-ray emission being originated by AGN photoionization, the RGS spectra studied by Gonzalez-Martin et al. (2010b) showed that in at least 30 per cent of their sample, a contribution of photoionization by the AGNs is required due to the presence of radiative recombination continua (RRC) from CV emission line. However, this does not guarantee a dominance of this emission mechanism. Moreover, cone-like morphologies at soft X-rays in some objects in this study point out again to the photoionization by the AGNs being responsible for the soft X-ray emission. Nevertheless, this assumes that we are seing LINERs with an LOS perpendicular to the accretion disc, which might not be the case. Indeed, the ultraviolet and X-ray variability detected for many of these LINERs (Maoz et al. 2005; Hernández-García et al. 2015) is in favour of a direct view of the AGNs (i.e. perpendicular to the disc under the UM). This is consistent with the fact that most of the [O III] morphologies found for LINERs are spheroids, if we assume that the [O III] traces the NLR. In addition, the fact that we detected a clear correspondence between soft X-ray and [O III] morphologies only in objects with log (LHX)>40, and also that all the objects where soft X-rays and [O III] match their morphologies seem to better follow the previously found relation between the size of the region and the hard X-ray luminosity (see Fig. 2 and Section 5.2), may argue in favour of the scenario in which the AGNs do not have enough thrust to ionize in the low-luminosity regime (Elitzur & Shlosman 2006; Elitzur & Ho 2009), ruling out photoionization by the AGNs at both soft X-ray and [O III] emissions. In this case, the most reasonable explanation for the [O III] is the host galaxy emission, which, anyhow, could also be on top of the AGNs, preventing its detection and erasing the connection (González-Martín et al. 2014). The host galaxy can contribute either as star formation or shocks to the total [O III] emission. Regarding the soft X-ray origin, Mingo et al. (2014) confirmed that jets are the main responsible for soft X-ray emission from their sources. In our sample, jets are identified in NGC 1052 (Kadler et al. 2004), where the jet position angle would be consistent with the extended soft X-ray emission shown here.
Citation: Hernández-García et al. 2015
Citation function: Similarities; Similarities
Citation: Maoz et al. 2005
Citation function: Compare/Contrast; Compare/Contrast; Similarities
A huge amount of data at all wavelengths has been collected in the past 20 years on black hole X-ray binaries, hereafter XrBs (for a global review see Dunn et al. 2010). These objects spend most of their time in quiescence at very low accretion rates, but occasionally, they produce outbursts that last from a few months to a year. Their flux then rises by several orders of magnitude across the whole electromagnetic spectrum (see, e.g., Corbel et al. 2004, Fender et al. 2006; Remillard & McClintock 2006; Done et al. 2007, for recent reviews). During an outburst, XrBs show very different spectral and temporal states that can be easily distinguished in a hardness-intensity diagram (HID) where the X-ray luminosity is plotted versus the hardness ratio of the X-ray spectrum (see, e.g., Homan et al. 2001; Fender et al. 2004). The evolutionary track produces a typical q-shaped figure that reveals a hysteresis: outbursting XrBs have two distinct spectra with the same X-ray luminosity above 1–2% Eddington luminosity. At the beginning of the outburst, the system is in the so-called hard state: the spectrum has a hard power-law shape up to a few tens to hundreds of keV, requiring a very hot, optically thin plasma (referred to as the “corona”). Then, when the system reaches high luminosities (up to a few tens of the Eddington luminosity), it transits within a few days through a bright intermediate state into the so-called soft state (referred to as the “cold disk”). In this state, the spectrum is dominated by strong and soft X-ray emission, which is commonly interpreted as thermal emission from an optically thick geometrically thin accretion flow. In the latter state, the luminosity starts to decrease and the system returns to the hard state, transiting through a faint intermediate state. The luminosities at which a system transits from hard to soft states are several times higher than the luminosity of the reverse transition (see Appendix in Dunn et al. 2010).
Citation: Remillard & McClintock 2006
Citation function: Background
EE Cet (ADS 2163 B) is the southern (slightly fainter) component of the visual binary WDS 02499+0856 (Mason et al. 2001). It was discovered by the HIPPARCOS mission (Perryman et al. 1997), by noticing the variability of the combined light of both visual components. Lampens et al. (2001) performed photometric measurements of the visual pair and gave (but only for one epoch) the following values V(A) = 9.47 mag and V(B) = 9.83 mag. Pribulla & Rucinski (2006) lists the orbital parameters (orientation and separation) of WDS 02499+0856 and gave θ = 194°, ρ = 5.66 arcsec and magnitude difference ΔV = 0.07 mag (the magnitude difference can be as large as ΔV = 0.36 mag, due to photometric variability of the eclipsing binary). WDS 02499+0856 turned out to be a quadruple system, when the northern component was found to be a double-lined (SB2) binary from the DDO spectroscopic observations (Pribulla & Rucinski 2006). D’Angelo et al. (2006) re-confirmed the multiplicity of the system, and listed it among the contact binaries with additional components. Radial velocity observations from Rucinski et al. (2002) resulted in a well-defined circular orbit of the contact binary, with K1 = 84.05 km s−1, K2 = 266.92 km s−1 (q = 0.315), and an F8V spectral type. Karami & Mohebi (2007) using their own velocity curve analysis method, arrived at almost identical results for the mass ratio. Djurašević et al. (2006) presented the first model, resulting in orbital inclination of i = 78.5° and a fill-out factor of f = 32.69 per cent, T2 = 6314 K, and T1 = 6095 K, when spots were added. Their no-spot model resulted in very close value for the fill-out factor but slightly different geometrical and orbital parameters. The physical parameters derived in this study were: M1 = 1.37 M⊙, M2 = 0.43 M⊙, and mean radii R1 = 1.35 R⊙, R2 = 0.82 R⊙. It is worth noting here that the light curves analysed by these authors included the visual component in the photometric aperture with a contamination of about 54 per cent.
Citation: Pribulla & Rucinski (2006)
Citation function: Background
Citation: Pribulla & Rucinski 2006
Citation function: Background
SEP events with gigaelectronvolt particles are generally rare. Typically about a dozen events occur during each solar cycle, although only two GLEs were reported in cycle 24, probably due to the change in the state of the heliosphere (Gopalswamy et al. 2013a, 2014a; Thakur et al. 2014). It appears that the 2012 July 23 event would have been another GLE event if it had occurred on the front side of the Sun. The purpose of this paper is to examine the event from the perspectives of CME kinematics, SEP intensity and spectrum, and radio-burst association to see if the 2012 July 23 event can be considered as an extreme particle event. The reason for considering these properties is clear from the following facts. Particles up to gigaelectronvolt energies are accelerated by strong shocks driven by CMEs of very high speeds (∼2000 km s−1) and intense, soft X-ray flares (see Gopalswamy et al. 2010, 2012b). The high speed is typically attained very close to the Sun, so the density and magnetic field in the corona are high for efficient particle acceleration (e.g., Mewaldt et al. 2012; Gopalswamy et al. 2014a). The high CME speed implies that a fast-mode MHD shock forms close to the Sun, as indicated by the onset of metric type II radio bursts, typically at heights 1.5 solar radii (Rs). CMEs attaining high speeds near the Sun have to accelerate impulsively, so these events are characterized by high initial acceleration (∼2 km s−2, see Gopalswamy et al. 2012b). This is in contrast to slowly accelerating CMEs (from filament regions outside active regions) that form shocks at large distances from the Sun and do not accelerate particles to energies more than a few tens of megaelectronvolts (Gopalswamy et al. 2015a, 2015d). Accordingly, the SEP spectra of such events are very soft, as opposed to the hard spectra of GLE events. Whether an event has a soft or hard spectrum is important information because the hard-spectrum events have stronger space weather impacts (see, e.g., Reames 2013). SEP events with gigaelectronvolt components are accompanied by type II radio bursts from meter (m) wavelengths to kilometer (km) wavelengths (Gopalswamy et al. 2005b, 2010). Type II bursts occurring at such wide-ranging wavelengths imply strong shocks throughout the inner heliosphere (Gopalswamy et al. 2005a).
Citation: Gopalswamy et al. 2010
Citation function: Background
Citation: Gopalswamy et al.; 2014a
Citation function: Background
Citation: Gopalswamy et al.; 2010
Citation function: Background
Vortex formation in protoplanetary discs has been shown to proceed through various mechanisms. Klahr & Bodenheimer (2003) demonstrated that in discs with a radial entropy gradient, azimuthal perturbations are baroclinically unstable and can lead to the formation of vortices. The most commonly invoked mechanism to form a large-scale vortex in a protoplanetary disc is the Rossby wave instability (RWI). The RWI can be triggered at steep radial density gradients (Lovelace et al. 1999; Li et al. 2000, 2001); such a gradient may be generated at the viscosity transition between regions in which the magnetorotational instability can operate (e.g. Varnière & Tagger 2006; Lyra et al. 2009b; Lyra & Mac Low 2012; Regály et al. 2012). Steep radial density gradients are also found at the outer edge of a gap carved in the disc by a giant planet. Indeed, hydrodynamics simulations of protoplanetary discs containing a giant planet have successfully produced large-scale vortices (e.g. Li et al. 2005; Masset et al. 2006; de Val-Borro et al. 2007; Meheut et al. 2010; Lin & Papaloizou 2011; Lin 2012), and Regály, Juhász & Nehéz (2017) presented a comparison of vortices formed at a gap edge to those formed at a viscosity transition. However, Hammer, Kratter & Lin (2017) pointed out that the short growth time-scales of the giant planet prescribed in these simulations result in unphysically large perturbations that are capable of setting up the RWI-unstable density gradient at the gap’s outer edge. When the giant planet is grown over more realistic time-scales, vortex production is significantly suppressed because viscosity has time to smooth out any steep gradients and the planetary torque itself can reshape the gap edge. Under the most favourable conditions for vortex formation, vortex lifetimes were found to be limited. One may then argue that if these asymmetries are indeed vortices formed through the RWI triggered by giant planets, we must be seeing them just after the planet has formed. However, if a RWI-generated vortex typically survives for ∼103 orbits and discs are ∼105 orbits old, observing the vortex so soon after its formation is unlikely. It is therefore uncertain whether the non-axisymmetric substructures seen in observations can be explained by this mechanism.
Citation: Regály et al. 2012
Citation function: Background; Background
Our initial sample comprised all 40 stars of the MMT-HVS survey that were revised by Kreuzer et al. (2020), together with 30 runaway stars from the collection of Silva & Napiwotzki (2011) for which we were able to obtain spectra. This covers the majority of objects with high ejection velocities in that compilation. This group was complemented by the prototype hyper-runaway star HD 271791 (Heber et al. 2008), the potential hyper-runaway stars SDSS J013655.91+242546.0 (J0136+2425 for short, Tillich et al. 2009) and HIP 60350 (Irrgang et al. 2010), the extreme disk-runaway star PG 1610+062 (Irrgang et al. 2019), the four candidate HVSs from the LAMOST survey (Zheng et al. 2014; Huang et al. 2017; Li et al. 2018), and the probable Hills star S5-HVS 1 (Koposov et al. 2020). Based on proper motions from Gaia EDR3, we then carried out spectroscopic and kinematic analyses (see Sects. 3 and 4 for details) of all members of this initial sample to filter out those with both a high ejection velocity and a relatively well-constrained origin within the Galactic disk. In order to account for the fact that massive stars are typically ejected at lower velocity, we chose a mass-dependent threshold for the deduced 1σ upper limit of the ejection velocity, that is, 400 km s−1 or 320 km s−1 for stars with masses below or above 5 M⊙, respectively. The first cut applies to almost all stars from the MMT-HVS survey, while most of the others fall into the second category. The chosen thresholds roughly represent the limits for the classical ejection scenarios (see, e.g., Tauris 2015; Irrgang et al. 2018a, and references therein). A disk origin was granted when the 1σ lower limit of the inferred galactocentric plane-crossing radius was below 25 kpc (motivated by Xu et al. 2015), while visual inspection was used to judge whether the origin was sufficiently well constrained. These criteria left us with 14 stars from the MMT-HVS survey, 7 stars from the Silva & Napiwotzki (2011) sample, all 4 stars from the LAMOST survey, and the 5 individual targets mentioned above, yielding a final sample of 30 program stars, the names of which are listed in Table 1.
Citation: Zheng et al. 2014
Citation function: Uses; Uses