adsabs/FOCAL huggingface dataset:
However, a significant outstanding problem of using red giants is that modeling their individual frequencies is too time consuming for the analysis of tens of thousands of stars. We therefore rely on using asteroseismic scaling relations, 






ν


max


∝


gT


eff


−
1

/

2




 and 




Δ
ν
∝
ρ


 (Brown et al. 1991; Kjeldsen & Bedding 1995), to estimate their radius and mass (and hence age). Here, 






ν


max




 is the frequency of maximum amplitude and 




Δ
ν


 the average large frequency separation. These relations assume that the structure of a red giant star is homologous with respect to the Sun. In reality this assumption is not strictly correct and verification of the relations is ongoing. However, the independent high-precision estimates of mass and radius required for this verification are difficult to obtain. For subgiants and dwarfs, the 






ν


max




 scaling relation has been shown to work well (Bedding 2014) and recently, Coelho et al. (2015) found the proportionality 






ν


max


∝


gT


eff


−
1

/

2




 to be accurate to within 1.5%. Using Hipparcos parallaxes and/or interferometry, the asteroseismic radii calculated from scaling relations have been found to be accurate to within 5% (Bruntt et al. 2010; Huber et al. 2012; Silva Aguirre et al. 2012). For giants we generally do not have accurate parallaxes, so such studies are awaiting results from Gaia (Perryman 2002). Open clusters have been used to test the scaling relations for giants (Brogaard et al. 2012; Miglio et al. 2012; Sandquist et al. 2013). Miglio et al. (2012) found agreement to within 5% for scaling relation-based radii. Testing of masses is more challenging. For a few cases where such verification have been performed, the scaling relation-based masses seem to be overestimated for giants (Miglio et al. 2012; Frandsen et al. 2013; Epstein et al. 2014). For two lower red giant branch stars (Epstein et al. 2014) find evidence that the mass estimated by using only 




Δ
ν


 (but with additional 






ν


max




 independent quantities) is lower compared to using both 




Δ
ν


 and 






ν


max




. Based on this they suggested that a modification to the 






ν


max




 scaling relation might be required. Theoretical modeling has suggested corrections to the 




Δ
ν


 scaling relation (Stello et al. 2009; White et al. 2011; Miglio et al. 2013a), but there has been no comprehensive study to verify the corrections. In relation to 






ν


max




, Houdek et al. (1999) and Chaplin et al. (2008) had suggested theoretically that 






ν


max




 coincides with the plateau of the damping rate with frequency. Belkacem et al. (2011) confirmed this for the Sun using SoHO GOLF observations. Balmforth (1992) suggested that this is caused by a resonance between the thermal adjustment time of the superadiabatic boundary layer and the mode frequency, which was also confirmed by the theoretical study of Belkacem et al. (2011; see also Belkacem 2012; Belkacem et al. 2013). However, there is currently no way to accurately predict 






ν


max




 from theory.
Citation: Bruntt et al. 2010
Citation function: Background
We wish to add to the work of Zhou et al. (2014) by deriving Ṁacc for LkCa15 b. To this end we use the upper limit on the flux from the Hα, Paγ and Paβ lines to estimate the accretion luminosity (Lacc) in each line. The equation Ṁacc = 1.25 (LaccR∗)/(GM∗) is used (Gullbring et al. 1998; Hartmann et al. 1998) and the results of Alcalá et al. (2014) are used to convert the luminosity of each line (Lline) to Lacc. Rpl is taken as 0.22 R⊙ and 0.34 R⊙ for masses of 6 MJup  and 15 MJup respectively. Rpl was estimated using the SETTL evolutionary models for 1 Myr (Allard et al. 2003). Our method is analogous to studies which use various accretion tracers to derive Ṁacc in low-mass stars and BDs (Alcalá et al. 2014; Whelan et al. 2014a), and therefore it is assumed here that the correlation between Lline and Ṁacc for low-mass pre-main sequence stars is applicable for planetary mass objects. Studies of Ṁacc using Lacc estimated from Lline have shown that the optimum way to estimate Ṁacc is to use several accretion tracers found in different wavelength regimes (Rigliaco et al. 2011, 2012). In this way any spread in Ṁacc due to different tracers probing different regimes of accretion or having a jet/wind component can be overcome. This is our reason for using more than one accretion tracer here. However, as many of the accretion tracers found in the spectrum of LkCa 15 show strong variability (refer to Fig. 2) we chose to limit our calculations to the lines with the least variability namely Hα, Paγ and Paβ. Here we are specifically referring to the presence of variable absorption features. For example in E1 the He I 1.083 μm line is double peaked but in E2 and E3 it is single peaked with a deep red-shifted absorption. The Hγ and Hβ lines also show red-shifted absorption in E2. Additionally, Hα is the line in which we are most likely to see a contribution from the planetary companion Close et al. (2014). It is found that for the 3σ upper limit on the flux, log (Ṁacc [ Hα ] ) = −9.0 to −9.2, log (Ṁacc [ Paγ ] ) = −9.1 to −9.3 and log (Ṁacc [ Paβ ] ) = −8.9 to −9.1. The range in Ṁacc for the individual lines comes from the estimated range in the mass of the companion. These estimates are in agreement with predictions for mass accretion rates from circumplanetary disks, which place Ṁacc for a 1−10 MJup planet in the range 10-9 to 10-8+M⊙/yr (Zhu 2015).
Citation: Allard et al. 2003
Citation function: Uses
The equations are evolved from 1 to 550 Myr for models M1 and M2 and from 13 to 550 Myr for model M3. All the results are compared to data of the ONC (Davies et al. 2014; Cieza & Baliber 2007), Cyg OB2 (Roquette et al. 2017), Upper Sco (Rebull et al. 2018), h Per (Moraux et al. 2013), the Pleiades (Rebull et al. 2016; Lodieu et al. 2019) and M37 (Hartman et al. 2009). The selection criteria applied to the observational data are shown in Table 2. The Cyg OB2 sample only contains stars with periods longer than 2 days because they are free from contaminants (for more details, see Roquette et al. 2017). For the Pleiades, we have cross matched the samples from Lodieu et al. (2019) and Rebull et al. (2016). In this work, we adopt solar metallicity stellar evolutionary models (Baraffe et al. 1998). Most of the samples analyzed here have solar metallicity. The ONC (Padgett 1996; D’Orazi et al. 2009; Biazzo, Randich & Palla 2011), U Sco (Viana Almeida et al. 2009) and the Pleiades (Takeda, Hashimoto & Honda 2017) present solar metallicity. Although Cyg OB2 does not have determinations of abundance, Wright et al. (2010), Guarcello et al. (2013) and Roquette et al. (2017) all use solar metallicity stellar models. Berlanas et al. (2018) find no evidence of self-enrichment in Cyg OB2. There are some works showing that the metallicity of h Per is sub solar (Z  = 0.01 - Southworth, Maxted & Smalley 2004; Tamajo, Pavlovski & Southworth 2011). However, Dufton et al. (1990) and Smartt & Rolleston (1997) obtained solar abundances for the same cluster. For M37 (NGC 2099), Casamiquela et al. (2017) obtained Z  = + 0.08 while Netopil et al. (2016) obtained Z  = + 0.02. There are several uncertainties related to age determination particularly for young clusters (see, for example Soderblom et al. 2014). The ages and corresponding references adopted in this work are shown in the 2nd and 3rd columns of Table 2. For NGC 2362, Mayne & Naylor (2008) suggest an age range between 4 and 5 Myr. We choose the age of 5 Myr because this is in better agreement with our models. For Cyg OB2, Wright, Drew & Mohr-Smith (2015) observe an age spread consistent with an extended stellar formation phase of 6 Myr with a peak between 4 and 5 Myr. We assume an age of 5 Myr for the cluster, as suggested by Wright et al. (2010). Among the clusters studied in this work, the age attributed to U Sco exhibits the greatest variation, from 5 Myr (Slesnick, Hillenbrand & Carpenter 2008; Herczeg & Hillenbrand 2015) to 11 Myr (Pecaut et al. 2012). Rebull et al. (2018) adopt 8 Myr for it but an 11 Myr old rotational period distribution returns a better value in the KS tests performed in this work.
Citation: Takeda, Hashimoto & Honda 2017
Citation function: Uses; Uses
Magnetars are possibly the most exotic objects of all the classes of neutron stars. High magnetic fields (≳1014 G) and intense X-ray/soft γ-ray emission combined with erratic spin-down behaviour makes them unique amongst the neutron star population. The soft gamma-ray repeater (SGR) J1935+2154 was discovered when the Swift-Burst Alert Telescope detected an X-ray burst on 2014 July 5 (Lien et al. 2014; Stamatikos et al. 2014). This initial burst was then followed by a number of other short bursts (Cummings 2014). Following this, the sky location of SGR J1935+2154 was quickly recognized to place it coincident with the supernova remnant (SNR) G57.2+0.8, strongly suggesting an association (Gaensler 2014). Follow-up radio observations with the Very Large Array (VLA) at 6 GHz revealed no persistent radio emission down to 41 $\mathrm{\mu Jy}$ (3σ; Fong & Berger 2014). Subsequent observations with the Chandra X-ray telescope detected X-ray pulsations with a period of 3.24 s (Israel et al. 2014). Searches with the Giant Metrewave Radio Telescope (GMRT), Ooty, and Parkes telescopes failed to detect any pulsed radio emission (Burgay et al. 2014; Surnis et al. 2014). Since the initial X-ray outburst the source has undergone further outbursts in 2015 February, 2016 May, and 2019 November. It is noteworthy that sporadic X-ray bursts are seen (Kozlova et al. 2016; Younes et al. 2017a) even between these periods of activity, making it one of the most active magnetars known (Lin et al. 2020b). As well as the single pulse and periodicity searches, searches for continuum radio emission, which might be associated with a new pulsar wind nebula activated by the bursting activity, were carried out just after the discovery of the source, with the GMRT (Surnis et al. 2016). They report 3σ upper limits of 1.2 and 4.5 mJy at frequencies of 362.5 and 610 MHz, respectively. Searches for continuum emission associated with this most recent burst have been undertaken by a wide range of telescopes and we summarize these in Table 1. In spite of deep searches, no continuum emission has been detected at the best known position of SGR J1935+2154. A detection of pulsar-wind nebula emission from the magnetar would have significant implications on the emission physics of the magnetar and its association with G57.2+0.8, although a pulsar-wind nebula has been detected around only one magnetar to date (Younes et al. 2016).
Citation: Burgay et al. 2014
Citation function: Background
The extended star-forming region S235 is known to be a part of the giant molecular cloud G174+2.5 in the Perseus Spiral Arm (e.g., Heyer et al. 1996) and contains two known sites: the S235 complex and the S235AB region (see Figure 1 in Dewangan & Anandarao 2011 and also Figure 1 in Kirsanova et al. 2014). The present work is focused on the S235 complex and does not include the S235AB region. Different values of the distance (1.36, 1.59, 1.8, 2.1, and 2.5 kpc) to the extended star-forming region S235 have been reported in the literature (e.g., Georgelin et al. 1973; Israel & Felli 1978; Evans & Blair 1981; Brand & Blitz 1993; Burns et al. 2015; Foster & Brunt 2015). In the present work, we have chosen a distance of 1.8 kpc following Evans & Blair (1981), which is an intermediate value of the published distance range. The H ii region associated with the S235 complex is predominantly ionized by a single massive star, BD+35o1201, of O9.5V type (Georgelin et al. 1973). The S235 complex has been studied using multiple data sets spanning near-infrared (NIR) to radio wavelengths. The S235 complex is known as an active site of star formation, harboring young stellar clusters (e.g., Kirsanova et al. 2008; Camargo et al. 2011; Dewangan & Anandarao 2011; Chavarría et al. 2014) associated with known star-forming subregions, namely East 1, East 2, and the Central region (e.g., Kirsanova et al. 2008). In our previous work on the S235 complex using Spitzer-IRAC data (Dewangan & Anandarao 2011, hereafter Paper I) we detected several young stellar objects (YSOs) including a high mass protostellar object (HMPO) candidate as well as signatures of outflow activities. Using 13CO (1−0) line data, Kirsanova et al. (2008) found three molecular gas components (i.e., −18 km s−1





 km s−1 (red), −21 km s−1 





 km s−1 (central), and −25 km s−1 





 km s−1 (blue)) in the direction of the S235 complex. However, the complex is well traced in mainly two molecular gas components (central and blue). More recently, Kirsanova et al. (2014) derived physical parameters of dense gas (i.e., gas density and temperature) in subregions of the complex using ammonia (NH3) line observations. However, the properties of dense gas have not been explored with respect to the ionizing star location. Previous studies indicated that the S235 H ii region is interacting with its surrounding molecular cloud and the S235 complex has been cited as a possible site of triggered star formation (Kirsanova et al. 2008, 2014; Camargo et al. 2011).
Citation: Burns et al. 2015
Citation function: Compare/Contrast
In Fig. 7 we compare UGC 1378’s SFR density versus gas surface density (the Schmidt–Kennicutt relation) to data in the literature. The gas surface density corresponds to H i calculated from the 0th moment map from Mishra et al. (2017) in the same areas as SFR density. Points for the HSB and LSB discs are plotted as black and grey circles, respectively. We plot the mean SFR and H i surface density for the entire galaxy with a large open circle. The black line corresponds to the relation with an exponent of 1.4 found by Kennicutt (1998). Triangles give results for LSB galaxies published by Wyder et al. (2009), and bright and faint crosses show normal spiral galaxies from Kennicutt (1998) – total and H isurface densities. A blue line shows the best-fitting relation for the H isurface density of Bluedisk galaxies from Roychowdhury et al. (2015). We also plot the SFR in the outer regions of spiral galaxies (Bigiel et al. 2010, square symbols). In Fig. 7 the UGC 1378 measurements lie between normal spirals and LSB galaxies. The HSB disc data lie above the relation plotted for normal spirals, possibly indicating that the SFR is boosted by the bar driving gas to the star-forming rings. We cannot account for molecular gas since there are no available measurements for UGC 1378. The contribution of molecular gas would likely move the HSB disc of UGC 1378 towards the locus of normal galaxies. Because the HSB SFR of UGC 1378 is close to the predicted SFR from the Kennicutt (1998) relation obtained from H i densities (faint crosses in Fig. 7). The LSB disc of UGC 1378 lies below the correlation and accounting for molecular gas would only increase the deviation from the normal Schmidt–Kennicutt relation. Similar deviations are observed in ‘classical’ LSB galaxies, Bluedisk galaxies (Roychowdhury et al. 2015), outer parts of HSB spiral galaxies (Bigiel et al. 2010), and H idiscs in early-type galaxies (Yıldız et al. 2017). These deviations for LSB galaxies are at least partially explained by their lower gas densities leading to lower SFRs (Abramova & Zasov 2011). A recent episode of gas accretion on to the disc of UGC 1378 may also contribute to a lower SFR if the gas is not yet fully participating in the star formation. Lutz et al. (2017) studied a sample of very H i rich galaxies and proposed that very high specific angular momentum in H irich galaxies prevents the accreted gas from being transported to the mid-plane of the disc and being converted into stars. This mechanism may act to preserve giant gaseous discs.
Citation: Bigiel et al. 2010; Bigiel et al. 2010
Citation function: Uses; Similarities; Similarities
Applied to massive stars, one can expect the presence of a tertiary companion to enrich the variety of evolutionary pathways in the inner binary by driving it to close stellar interactions (Toonen, Hamers & Portegies Zwart 2016; Toonen et al. 2020). Yet, simulating the evolution of massive stellar triples poses a difficult challenge since the stellar physics of each individual star and the gravitational three-body dynamics have to be combined in a self-consistent way. Concerning massive stars, both of these aspects are closely intertwined. For instance, kicks experienced in the supernova (SN) explosions modify and potentially disrupt the three-body configuration. Likewise, massive stars at high metallicity suffer significant mass-loss through stellar winds that loosen the inner and outer orbits (Castor, Abbott & Klein 1975; Vanbeveren 1991; Vink, de Koter & Lamers 2001; Schneider et al. 2015; Mapelli 2016). It has been shown that mass-loss in the inner binary due to winds or at compact object formation could induce or strengthen the LK effect (Shappee & Thompson 2013; Michaely & Perets 2014). In addition, massive stars attain large radii as they evolve off the main sequence and beyond, so that Roche lobe overflow and mergers are expected to occur frequently in isolated massive binaries (Bonnell & Bate 2005; Eldridge, Izzard & Tout 2008; Sana et al. 2012; Schneider, Podsiadlowski & Müller 2021). For example, more than $\sim 70\, {{\ \rm per\ cent}}$ of Galactic massive O-type stars are expected to undergo at least one mass-transfer episode with their binary companion (Sana et al. 2012). A tertiary companion could facilitate these types of close stellar interaction via the LK mechanism by driving the inner binary to smaller pericentre distances. Previous studies of stellar triples have shown that these may give rise to X-ray binaries (Naoz et al. 2016) or even trigger a stellar merger (Perets & Kratter 2012; Stegmann et al. 2022) leading to the formation of Blue stragglers (Perets & Fabrycky 2009; Naoz & Fabrycky 2014; Antonini et al. 2016) and type Ia SN (Iben & Tutukov 1999; Thompson 2011). Moreover, an expanded tertiary star could itself overflow its Roche lobe and initiate a mass transfer phase on to the inner binary (de Vries, Portegies Zwart & Figueira 2014; Portegies Zwart & Leigh 2019; Di Stefano 2020a, b; Hamers et al. 2021).
Citation: Mapelli 2016
Citation function: Background
Our free chemistry retrieval suite results are further at odds with equilibrium chemistry due to the potential presence of an overabundance of H− in many retrievals across our suite. Across our free chemistry retrieval suite, the increasing abundance of H− is driven by the downturn in the final four WFC3/G141 data points, as seen, for example, in retrievals A and C in Figs 11 and 12. Here, H− acts like a uniform opacity source in the optical, which turns off at around 1.6 $\rm{\mu m}$ coinciding with these data points. Under equilibrium chemistry conditions, the abundances of H− retrieved in our suite would be possible only at temperatures ≳2500 K (Kitzmann & Heng 2018). However, the production of appreciable quantities of H− could be possible under disequilibrium processes. As described by Lewis et al. (2020), enhanced e-, H, and H2 mixing ratios (Lavvas, Koskinen & Yelle 2014) with production of H− by H2 dissociative electron attachment and destruction by atomic H collisional detachment can enable mixing ratios of H−, which are at the orders of magnitude of the retrieved abundances, even for the temperatures expected within the atmosphere of WASP-17b. The conditions required for these processes are likely for hot Jupiters orbiting F-type stars such as WASP-17, making the abundances of H− seen in our retrieval suite plausible outside of equilibrium conditions. While the downturn at 1.6 $\rm{\mu m}$, which seems to drive the inclusion of H−, could be explained by patchy clouds (e.g. Line & Parmentier 2016; MacDonald & Madhusudhan 2017), our explorations of the cloud parametrization found patchy clouds to be statistically disfavoured for the measured transmission spectrum (see Section 6.1). Additional IR observations around 1.6 $\rm{\mu m}$, which overlap with the existing G141 data, would help to corroborate the shape of the transmission spectrum in this region, shedding light as to whether H− is indeed present within the atmosphere of WASP-17b.
Citation: Kitzmann & Heng 2018
Citation function: Differences; Differences
Theoretically, the production, propagation, and escape of LyC photons are related to the physical properties of the galaxies. Firstly, the production of LyC radiation implies the presence of young, massive stars, and therefore of on-going star formation. Because of the fast recombination timescale of the HI atoms, previous episodes of star formation have no significant impact on the production of ionizing photons eventually escaping the galaxy (e.g. Paardekooper et al. 2015). Secondly, the propagation of LyC photons within the ISM is favoured by a negligible amount of dust and low column density of HI (N(HI) ≤ 1018 cm-2) in a 10-pc scale region around the emitting star clusters. This could be the case for galaxies embedded in dark-matter halos with masses less than 108 M⊙ (Yajima et al. 2011; Wise et al. 2014; Paardekooper et al. 2015). However, even galaxies residing in more massive halos can have lines of sight favourable to the propagation and the escape of LyC photons (Gnedin et al. 2008; Roy et al. 2015). Supernova explosions could have cleared their ISM, and star-formation episodes could occur in their outskirts. In addition, “runaway” OB stars up to 1 kpc away from the initial-origin regions are proposed to significantly contribute to the amount of LyC photons finally emitted into the IGM (Conroy & Kratter 2012). Thirdly, LyC photons emitted into the IGM can affect the galaxy environment, changing the ratio of neutral vs ionized gas, eventually fuelling the ISM (e.g. Martin et al. 2012, for a study of ionized-metal outflows and inflows). Simulations at intermediate redshift have shown that the LyC escape fraction (fesc(LyC)) steeply decreases as the dark-matter halo mass (Mh) increases at 3 z 6 (e.g. Yajima et al. 2011) and that the median fesc(LyC) also changes with redshift at z = 4−6 (e.g., Cen & Kimm 2015). It is worth stressing that while some authors find that the LyC escape fraction decreases with the increase in the halo mass (see also Ferrara & Loeb 2013), other works find the opposite trend: fesc(LyC) is found to range from a few percent (e.g. Gnedin et al. 2008) up to 
        20−30% (e.g. Mitra et al. 2013) or even higher (e.g. Wise & Cen 2009).
Citation: Roy et al. 2015
Citation function: Background