adsabs/FOCAL huggingface dataset:
Based on indirect evidence, there are two main candidates for the γ-ray emitting region. The first one is close to the BH and the γ-ray emission is produced inside the broad-line region (BLR; e.g., Dermer & Schlickeiser 1993; Blandford & Levinson 1995; Ghisellini & Madau 1996; Georganopoulos et al. 2001; Fan et al. 2006; Bai et al. 2009; Tavecchio & Mazin 2009; Isler et al. 2013; Hu et al. 2015). This argument is supported by constraining the size of the γ-ray emission region through the magnification factor during the γ-ray variability in some gravitationally lensed γ-ray blazars (e.g., Neronov et al. 2015; Vovk & Neronov 2016), the short variability timescales (down to several hours; e.g., Tavecchio et al. 2010), and the sharp breaks at the GeV band seen in the γ-ray spectra of some FSRQs that may caused by the opacity to pair production (e.g., Liu & Bai 2006; Bai et al. 2009; Poutanen & Stern 2010; Stern & Poutanen 2011). The second possible region for the γ-ray production is beyond the BLR (e.g., Błażejowski et al. 2000; Arbeiter et al. 2002; Sokolov & Marscher 2005; Yan et al. 2012; Meyer et al. 2015). Some multi-wavelength observations for the giant flares suggested that both the sub-millimeter and γ-rays are produced 10–20 pc from the BH (e.g., Larionov et al. 2008; Sikora et al. 2008), where it was assumed that γ-ray and radio emission is triggered by shocks propagating along a relativistic jet. The time delay between flares in radio and γ-ray bands combined with the VLBI high-resolution observation on the size of radio core also indicated that the γ-ray emitting region stays in the upstream of, but not far from the radio core (e.g., Jorstad et al. 2010; Marscher & Jorstad 2010; Pushkarev et al. 2010; Agudo et al. 2011b; Fuhrmann et al. 2014; Max-Moerbeck et al. 2014; Ramakrishnan et al. 2016). It is normally believed that γ-ray emission in blazars come from the inverse-Compton (IC) process, where external Compton (EC) plays an important role in LSP blazars, while ISP and HSP blazars (BL Lacs) generally can be explained by the synchrotron self-Compton (SSC; e.g., Krawczynski et al. 2004; Chen & Bai 2011; Kang et al. 2014; Zhang et al. 2014). In the EC process, the seed photons are determined by the location of the γ-ray emitting region, which may be dominated by an accretion disk, BLR, infrared torus, and cosmic background, respectively, if the γ-ray emitting region is located near the BH horizon, inside the BLR, inside the torus, or much beyond of the torus (e.g., Ghisellini & Tavecchio 2009). Kang et al. (2014) compared the seed photons from the BLR and dusty torus in the SED fitting of LSP blazars, and found that the seed photons from torus are preferred based on a 





 method. Zheng et al. (2017) also explored this issue using a stratified jet and found that the seed photons from torus provide a better match to the observations.
Citation: Stern & Poutanen 2011
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The study of cometary plasma composition has been subjected to a great interest after the ion mass spectrometer onboard Giotto spacecraft detected many peaks in the mass range 12 and 120 amu (Balsiger et al. 1986; Krankowsky et al. 1986; Mitchell et al. 1987; Altwegg et al. 1993). By developing photochemical models, numerous studies focused on comet 1P/Halley explained the observed ion distribution in a water-dominated coma (Allen et al. 1987; Wegmann et al. 1987; Schmidt et al. 1988; Cravens 1989; Bhardwaj, Haider & Singhal 1990, 1996; Gan & Cravens 1990; Ip et al. 1990; Haider, Bhardwaj & Singhal 1993; Häberli et al. 1995; Haider & Bhardwaj 1997, 2005; Bhardwaj 1999; Rubin et al. 2009; Cordiner & Charnley 2014). By making 2 yr of observations, the recent Rosetta space mission on comet 67P/Churyumov–Gerasimenko has revolutionized our understanding of the activity of the cometary coma. During the Rosetta observation period, continuous measurements around the nucleus were helpful to study the evolution of ion and neutral distribution and also the driving photochemical processes in the coma. Several modelling works on this comet have shown that ion composition in the coma varies based on the sublimation rate of the nucleus (Vigren & Galand 2013; Fuselier et al. 2015, 2016; Galand et al. 2016; Heritier et al. 2017, 2018; Vigren et al. 2017; Beth, Galand & Heritier 2019). All these studies show that solar photons are the primary energy source that determines the ion composition in the inner coma. Solar extreme ultraviolet photons having an energy more than 12 eV ionize H2O and produce H2O+, and the collisions among these species quickly lead to the formation of H3O+. The sublimated parent species such as CH3OH, NH3, HCN, HCOOH, and CH3CHO, have high proton affinities compared to that of H2O, causing the loss of H3O+ in the inner coma. Haider & Bhardwaj (2005) developed a comprehensive chemical network to study the ion distribution in comet 1P/Halley. Their calculations show that NH$_4^+$ is the most dominant ion in the inner coma followed by H3O+ and CH3OH$_2^+$ ions. Similarly, the model calculations of Heritier et al. (2017) on comet 67P/Churyumov–Gerasimenko showed that NH$_4^+$, CH3OH$_2^+$, H3O+, H3S+, and HCNH+ are the important ions in the inner coma. They also showed that the densities of these ions vary with the relative mixing ratios of corresponding proton affinity species coming from the nucleus. Even if the mixing ratios of parent species, which have high proton affinity, are very low (2 per cent), they can play a significant role in modifying the ionospheric composition of the inner coma. Hence, the ion distribution in the cometary coma essentially depends on the neutral composition and photochemical reactions.
Citation: Rubin et al. 2009
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The main challenge in detecting the cosmological H i signal, common to all of these experiments, is the strong contamination of systematic effects (ionospheric distortion, telescope response, calibration, etc.) and bright foregrounds (Galactic and extragalactic) (Datta, Bhatnagar & Carilli 2009). Foreground sources include diffuse Galactic synchrotron emission (DGSE) from our Galaxy (Shaver et al. 1999), free–free emission from Galactic and extragalactic sources (Cooray & Furlanetto 2004), faint radio-loud quasars (Di Matteo et al. 2002), synchrotron emission from low-redshift Galaxy clusters (Di Matteo, Ciardi & Miniati 2004), extragalactic point sources, etc. Typically, foregrounds are four to five orders of magnitude stronger than the redshifted H i signal (Zaldarriaga, Furlanetto & Hernquist 2004; Bharadwaj & Ali 2005; Jelić et al. 2008; Bernardi et al. 2009; Jelić et al. 2010; Zaroubi et al. 2012; Chapman et al. 2015). There are several different ways to deal with foregrounds, but all the methods rely on the fact that foreground sources have a smooth spectral shape. However, the redshifted H i 21 cm signal has spectral structure (Pritchard & Loeb 2012). In fact, this difference in spectral properties between the strong foreground and faint 21 cm signal can be used favourably for the detection of the cosmological signal (Datta, Bowman & Carilli 2010). Hence, accurate knowledge of the spectral ‘smoothness’ of the foreground becomes critical. Our current study makes an attempt to constrain the spectral behaviour of the foregrounds near the redshifted 21 cm signal frequencies. The three main techniques proposed to overcome foreground contamination are foreground avoidance, foreground removal and foreground suppression. Instead of an isotropic 1D power spectrum, P(k), of H i brightness temperature fluctuation, a cylindrical (2D) power spectrum, P(k⊥, k∥) is a useful diagnostic in terms of foreground avoidance. Spectral smoothness of foregrounds confines the majority of foreground power to low k∥ modes, resulting in ‘foreground wedge’. In the foreground avoidance technique, the EoR signal is searched for outside this wedge, in the so-called ‘EoR window’ (Datta et al. 2010; Parsons et al. 2012; Vedantham, Udaya Shankar & Subrahmanyan 2012; Pober et al. 2013a; Thyagarajan et al. 2013; Dillon et al. 2015). However, errors in calibration of chromatic instruments and insufficient knowledge of the wedge boundary can leak foreground power into the wedge and consequently detection of the 21 cm signal becomes challenging even inside the EoR window. Foregrounds can be modelled very precisely and subtracted from the data set. Also, without any modelling a component analysis method can be used to mitigate foregrounds (for details see Chapman et al. 2012, 2013). The foregrounds can be suppressed by weighting foreground-dominated modes appropriately (Liu & Tegmark 2011).
Citation: Jelić et al. 2010
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Citation: Zaldarriaga, Furlanetto & Hernquist 2004
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The oEA stars have pulsation features similar to classical δ Sct stars. However, their pulsations may be influenced by the tidal interaction and mass transfer between the components, as well as gravitational force from companions. Some pulsations in KIC 6220497 can be excited by the tidal forces of the secondary companion. Tidally excited modes occur when the orbital frequency is close to a stellar eigenfrequency in a binary star with an eccentric orbit. The signature of the pulsation modes is the frequencies at integer multiples of the orbital frequency (Welsh et al. 2011; Thompson et al. 2012; Hambleton et al. 2013). We detected four (f6, f11, f20, f24) frequencies that are the harmonics of the orbital frequency, which could be partly affected by imperfect removal of the eclipses from the light curve. Although the orbit of the binary system is circular, these frequencies could result from tidally induced pulsations (Reyniers & Smeyers 2003a,b; Southworth et al. 2011). In the upper panel of Fig. 8, we plot the dominant pulsation period versus the orbital period for 74 oEA stars, including KIC 6220497. The data are taken from the compilations of Zhang et al. (2013, 67 oEA stars) and from more recent literature (Yang, Wei & Li 2014 for FR Ori; Zhang et al. 2014 for OO Dra; Zhang, Luo & Wang 2015a for EW Boo; Zhang et al. 2015b for V392 Ori; Lee et al. 2016 for KIC 4739791; Soydugan et al. 2016 for XZ Aql). In the panel, we can see that the pulsation period of KIC 6220497 deviated from the general trend of the oEA stars and also the empirical relation between Ppul and Porb of Zhang et al. (2013). The pulsation periods in EBs increase with decreasing gravitational pull exerted by the secondary companion to the pulsating primary component (Soydugan et al. 2006). For KIC 6220497, the gravitational pull applied to per gram of the matter on the surface of the pulsating component by the lobe-filling secondary was calculated to be log F = 2.90 in cgs units, which is about 3.4 times smaller than the value of 9.89 (again, in cgs units) taken from the relation of log Ppul = −0.61log F + 5.1 calibrated by Soydugan et al. (2006). However, as shown in the lower panels of Fig. 8 from 34 oEA stars currently known, the gravitational force log F for KIC 6220497 matches well with those for the other oEA stars with similar orbital periods, while the surface gravity log g1 = 3.78 for the primary component is clearly smaller. The period–gravity relation is similar to that for radially pulsating stars suggested by Fernie (1995); as the surface gravity decreases, its pulsation period increases. Further, it is known that the more evolved the star, the slower the pulsations (cf. Liakos et al. 2012). The pulsation period and the surface gravity of KIC 6220497 indicate that the system might be a more evolved EB than the other oEA stars. We think that the pulsation periods strongly depend on the surface gravities of the pulsating components and the evolutionary status of the binary stars.
Citation: Zhang, Luo & Wang 2015a
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As we describe in Section 2 the case of TNDWs propagating in a plasma with equal mass fraction of 12C and 16O (CO) and density of ρ0,7 ≈ 1,1 which is typical for Type Ia supernovae, is particularly challenging for full-star simulations. In addition to the problem that the burning length-scale is much smaller than the typical cell size, near detailed balance is obtained for many isotopes while NSE is not reached. We test in Section 2 two available one-dimensional (1D) codes: a modified version of the 1D, Lagrangian version of the vulcan code (hereafter V1D; for details, see Livne 1993) and a modified version of the Eulerian, 1D hydrodynamic flash4.0 code with thermonuclear burning (Fryxell et al. 2000; Dubey et al. 2009), against the ρ0,7 = 1 case. We show that with resolutions that are typical for multidimensional full-star simulations, the V1D and the flash results are not satisfactory (up to $50{{\ \rm per\ cent}}$ error in V1D and up to $20{{\ \rm per\ cent}}$ error in flash). We demonstrate in Section 3 the operation of a new numerical scheme for thermonuclear burning that can be implemented in multidimensional full-star simulations. The new scheme allows an accurate calculation of TNDWs in a consistent way (i.e. without pre-describing the position and/or the conditions behind the TNDW) with all thermonuclear burning taking place in situ (without post-processing) for an arbitrary reaction network with hundreds of isotopes. The new scheme contains two important ingredients: (1) a burning limiter (a variant of Kushnir et al. 2013), which guarantees that the thermodynamic variables and the composition are accurate for the resolved scales, while keeping the numerical thermodynamic trajectory for unresolved scales within some controlled error from the true thermodynamic trajectory, and (2) adaptive statistical equilibrium (ASE) burning, which groups isotopes that are in detailed balance into one effective isotope, where the ratio between the isotope abundances inside the group is given from equilibrium conditions (this is an extension of the earlier attempts of Hix et al. 2007; Parete-Koon & Hix 2008; Parete-Koon, Hix & Thielemann 2008, 2010).
Citation: Fryxell et al. 2000
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WASP-43b is the most heavily scrutinized phase curve, with four analyses of this data set already published (Stevenson et al. 2017; Mendonça et al. 2018; Morello et al. 2019; May & Stevenson 2020). Our phase curve semi-amplitude, eclipse depth, and radius are consistent with all of these works. The more contentious issue is that of the phase curve’s phase offset and nightside temperature. Stevenson et al. (2017) initially reported only a 2σ upper limit on the nightside temperature of 650 K, while all subsequent reanalyses (including ours) favour a significantly detectable nightside temperature of ∼800 K. As for the planet’s phase offset, Stevenson et al. (2017) and May & Stevenson (2020) favour a larger phase offset (21 ± 2 °E) than Mendonça et al. (2018) and Morello et al. (2019) (12 ± 3 °E and 11 ± 2 °E). May & Stevenson (2020) claimed that the differences between the retrieved phase offsets is the result of temporal binning which was not used by Stevenson et al. (2017) and May & Stevenson (2020) but was used by Mendonça et al. (2018), Morello et al. (2019), and this work. Fitting the temporally binned photometry for all 17 phase curves with each of our detector models already required more than 2000 CPU hours, and expanding this to unbinned photometry for all phase curve fits would require more than 125 000 CPU hours (or 434 d using our 12× multithreading computer) optimistically assuming all of detector models scaled linearly with the number of input data. However, we did try fitting just the WASP-43b unbinned phase curve with our preferred detector model (BLISS) and found that our phase offset and nightside temperature was unchanged. Including a linear slope in time also did not affect our phase offset or nightside temperature. Instead, we find that the phase offset inferred by our models depends on the choice of phase curve model, as our 4-parameter (v2) phase curve models are consistent with those of Stevenson et al. (2017) and May & Stevenson (2020), while our 2-parameter phase curve models (v1) are consistent with Mendonça et al. (2018) and Morello et al. (2019). Ultimately, we cannot decide between these two discrepant offsets as the ΔBIC between the two phase curve models for our preferred BLISS detector model is only 3.7 (insignificantly favouring the 20.4 ± 3.6 offset from the v2 model). For reference, Stevenson et al. (2014b) found phase offsets ranging from roughly −6 to 17 deg east in the Hubble/WFC3 bandpass.
Citation: Stevenson et al. (2017)
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Citation: Stevenson et al. 2017
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We obtained atmospheric parameters using the high-resolution spectra. For our candidates with early spectral types: CoRoT–34, CoRoT 310204242, and CoRoT 110660135 we derived the effective temperature Teff, surface gravity log (g)sp, and metallicity [Fe/H] using grids of synthetic spectra, based on ATLAS12 model atmospheres (Kurucz 1996). The most important elements were considered using a hybrid non-local thermal equilibrium (NLTE) approach with the DETAIL/SURFACE package (Giddings 1981). Our model spectra, as well as the fitting method are described in detail in Irrgang et al. (2014) and Heuser (2018). For all other candidates, the synthetic spectra were calculated in LTE using the ATLAS12 and SYNTHE codes (Kurucz 1993). To consider all sensitive features in the observed spectra, we performed global χ2 fits. In the cases where combined spectra were available from several instruments, we fitted them simultaneously using the same model grid taking different resolution and radial velocity into account. Regions that were not well reproduced were removed from the fit. This includes the cores of hydrogen lines, lines that are not included in our model spectra, as well as lines with uncertain atomic data. It is challenging to determine the surface gravity for F-type stars, because it is correlated with both the Teff and the abundance of calcium and magnesium. We assumed systemic uncertainties of 0.2 dex for log (g)sp, of 0.1 dex for [Fe/H], and 2  per cent in Teff, which were added in quadrature to the much smaller statistical uncertainties. The atmospheric parameters from spectroscopy are listed in Table 2. Second, we derived stellar age, mass, and radius by fitting the observed parameters to mesa evolutionary tracks (Choi et al. 2016) in the (Teff, $\log \, g$, [Fe/H]) plane as described by Irrgang et al. (2016). For CoRoT–34, 35, and 36, we used the stellar mass and radius as input for the light-curve fit, which we used to constrain log (g)lc from the light curve directly (For details, see description of the light curve fit below). We derived the final stellar age, mass, and radius as listed in Table 2 by repeating the mesa fit using log (g)lc as prior. The best-fitting evolutionary tracks are shown in Fig. 2.
Citation: Kurucz 1993
Citation function: Uses
At present two main processes have been put forward for triggering star formation at the peripheries of H ii regions: collect-and-collapse (CC) and radiation-driven implosion (RDI; Elmegreen & Lada 1977; Deharveng et al. 2010). In the CC process, a compressed layer of neutral material is accumulated between the ionization front and shock front, and star formation occurs when this layer becomes gravitationally unstable. Unlike the CC process, in the RDI process a pre-existing denser gas is compressed by the shocks, and then forms stars. The CC process is more attractive because it allows the formation of massive stars or clusters (Deharveng et al. 2005). To investigate the star formation triggered by the CC process, several individual H ii regions have been studied, such as Sh 104 (Deharveng et al. 2003), RCW 79 (Zavagno et al. 2006), RCW 120 (Zavagno et al. 2007), Sh2-212 (Deharveng et al. 2008), Sh2-217 (Brand et al. 2011), Sh2-90 (Samal et al. 2014), Sh2-87 (Xu & Ju 2014), N6 (Yuan et al. 2014), and Gum 31 (Duronea et al. 2015). Moreover, there are also some statistical studies of infrared bubbles (e.g., Deharveng et al. 2010; Kendrew et al. 2012; Thompson et al. 2012; Kendrew et al. 2016), which are created by the expanding H ii regions. A common feature of the above studies is that several massive fragments are found on the border of each H ii region. Some star formation activity has been detected in these fragments, such as outflow, UCH ii, and water masers. Generally, if an expanding H ii region collects its surrounding gas, the morphology of the molecular gas will show three-dimensional spherical structure. CO observations can provide velocity information to reveal the gas structure surrounding the H ii regions. Beaumont & Williams (2010) observed 43 infrared bubbles using the CO molecular line, but they did not detect the front and back faces of these shells at blueshifted and redshifted velocities. Hence, they concluded that the bubbles enclosing H ii regions are two-dimensional rings formed in parental molecular clouds with thicknesses not greater than the bubble sizes. However, molecular gas with a three-dimensional spherical structure or a two-dimensional ring around H ii region is important for understanding the triggered star formation (Deharveng et al. 2015).
Citation: Zavagno et al. 2006
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One major difficulty afflicts the study of SNRs to learn about their supernovae: separating ejected material from swept-up surrounding unmodified interstellar medium (ISM) or modified CSM. Abundance clues are powerful but have limitations. Ideally, one would observe the youngest possible remnant that is large enough for adequate spatial resolution. That remnant appears to be the youngest Galactic SNR, G1.9+0.3 (Reynolds et al. 2008; see Figure 1). This object is about 



100′′

 in diameter, the smallest angular size of any confirmed Galactic SNR. Unfortunately, it is very highly absorbed, with an X-ray column density of about 5 × 1022 cm−2 (Reynolds et al. 2009), implying A

V
 ∼ 23
m
, so radio and X-rays are the only useful observational channels. The angular expansion rate of 0.64 arcsec yr−1 obtained from comparing X-ray images from 2007 and 2009 (Carlton et al. 2011) gives an upper limit for the age of about 160 yr, less if (as is almost certainly the case) deceleration has occurred; spatial variations in expansion rate (Borkowski et al. 2014) are consistent with an age of about 100 yr, or a date of around 1900. The high extinction would have rendered it unobservable in optical telescopes of that era. Furthermore, its X-ray spectrum is almost entirely synchrotron emission, making it a member of the small class of X-ray synchrotron–dominated SNRs. However, long observations with Chandra have allowed the detection of thermal emission from small regions (Borkowski et al. 2011, 2013b), with spectroscopic widths of ∼14,000 km s−1 confirming the large expansion proper motion, refined with a second Chandra observation (Carlton et al. 2011). The distance is still uncertain; the high column—higher than the entire Galactic column along nearby sight lines—suggests an association with the Galactic center, and a provisional distance of order 8.5 kpc has been assumed. Nearer would be very unlikely in view of the high absorption, but too much farther would make the expansion proper motion unreasonably large. An H i observation with the Giant Metrewave Radio Telescope (Roy & Pal 2014) has been used to set a lower limit of 10 kpc, certainly consistent with the known properties of G1.9+0.3.
Citation: Carlton et al. 2011
Citation function: Similarities; Similarities